{
    "0212/astro-ph0212491_arXiv.txt": {
        "abstract": " ",
        "introduction": "Since 1985~\\cite{bib-witten}, various experimental techniques have been proposed to detect the nuclear recoils that would be produced by the scattering of WIMP dark matter particles from our galactic halo (for a recent review of experimental searches, see e.g. Ref.~\\cite{bib-expreview}). The goal is to be sensitive to recoil energies down to 10 keV~\\cite{bib-lewin}, at event rates well below one per kilogram of detector material per day~\\cite{bib-rate}. For this, the detector response to low-energy nuclear recoils must be well known. Depending on the detection process (ionization, scintillation or heat), this response may differ significantly from that inferred from calibration with electron or gamma-ray sources. A quantity of interest is thus the {\\em quenching factor}, defined as the ratio of the signal amplitudes induced by a nuclear recoil and an electron of the same energy. This factor depends mainly on the detector material, the energy of the recoil and the detection process, although temperature and alignment effects may play a role (see e.g. Ref.~\\cite{bib-align}).  Quenching factor measurements have been made on scintillation detectors such as NaI~\\cite{bib-nai}, CsI~\\cite{bib-csi}, CaF$_2$~\\cite{bib-caf2} and Xe~\\cite{bib-xe}, leading to values ranging from 4 to 30\\% depending on the recoil nucleus and its energy. The measured quenching factors of ionization detectors such as Si~\\cite{bib-si} or Ge~\\cite{bib-chasman,bib-jones,bib-shutt,bib-messous} vary between 25 and 40\\% and appear to follow the recoil energy dependence predicted by the Lindhart theory~\\cite{bib-lindh}. More recently, the thermal detection efficiency for recoiling nuclei has been measured for the first time in a TeO$_{2}$ cryogenic detector, giving a thermal quenching factor slightly above unity\\cite{bib-aless}.  The purpose of the SICANE facility (SIte de CAlibration NEutron) is the measurement of quenching factors of cryogenic detectors using nuclear recoils induced by monoenergetic neutron beams. The facility has been commissioned with tests using scintillation in a NaI(Tl) crystal and measurements of quenching factors for ionization and heat signals in a germanium cryogenic detector developed by the EDELWEISS collaboration. ",
        "conclusions": "The commissioning tests of the SICANE array and its setup confirm its high relevance for the calibration of the response of cryogenic detectors to nuclear recoils.  Using inverse reactions to produce a highly collimated neutron beam, no massive shielding of the neutron detectors is necessary, giving more freedom in their spatial setup. Competitive results have been obtained in a short running time (less than a day) for the quenching of scintillation in NaI(Tl) and of ionization in Ge. The array also makes possible a deeper investigation of the effect of inelastic scattering on the measurement.  The powerful diagnostic tools provided by the array (simultaneous measurements at different angles, neutron identification and time-of-flight measurements) are ideally suited for the study of cryogenic detectors. For such measurements, important factors that require proper attention are: {\\em i)\\/} the slow response of the detector, and how it compares with the single rate in the cryogenic detector and the rate in coincidence with the neutron array; {\\em ii)\\/} the rate in the neutron detector array due to neutron scattering in the material in the vicinity of the detector and {\\em iii)\\/} the ability to resolve inelastic and elastic events, either using the energy or the time-of-flight measurements.  In addition, for the first time, it was directly verified that the ionization quenching factor of Ge at 35 mK is very close to that at liquid nitrogen temperature. This is consistent with the fact that the CDMS and EDELWEISS collaborations~\\cite{bib-shutt,bib-benoit} obtain $Q/Q'$ ratios compatible with the $Q$ values measured at liquid nitrogen temperature~\\cite{bib-chasman,bib-messous} under the assumption that $Q'$=1. Conversely, this last hypothesis is verified for the first time in the present work, albeit at a 15\\% level."
    },
    "0212/astro-ph0212172_arXiv.txt": {
        "abstract": "Millimeter and mid-infrared observations have been made of the dense clumps of dust and gas and of young stellar objects (YSOs)  associated with the bright, compact submillimeter source G79.3+0.3~P1 in the  relatively nearby MSX infrared-dark cloud  G79.3+0.3. The Gemini mid-infrared observations reported here indicate the presence of three YSOs within the cloud.  BIMA 3~mm continuum observations show that the brightest  of the YSOs is likely to be a Herbig Ae/Be star.  High-angular-resolution  molecular-line observations suggest that a wind from this star may be triggering collapse in the adjacent molecular cloud.  The submillimeter source G79.3+0.3~P1 itself does not contain infrared sources and may represent an earlier stage of star formation. ",
        "introduction": "\\citet{ega98} have identified a large population of cold dust clouds in images of the Galactic Plane made with the MSX satellite \\citep{pri01}.  These clouds appear as dark patches of  absorbing material against a background of mid-infrared emission bands and emission from warm, small dust grains.  Egan et al. defined the infrared-dark clouds (IRDCs) as having one to several magnitudes of extinction at 8~\\micron\\  and no obvious emission in any of the 8 to 25~\\micron\\ bands observed by the MSX satellite.  They concluded that the IRDCs possess hundreds of magnitudes of visual extinction and contain large column densities of cold dust. Observations of a selection of these clouds in H$_2$CO by \\citet{car98} revealed that much of the gas inside the IRDCs has T$_{\\rm K}$~$\\approx$~10-20~K and n(H$_2$)~$\\gtrsim$~10$^6$~cm$^{-3}$, with H$_2$ column densities ranging up to 10$^{23}$~cm$^{-2}$.  Most of these objects are quite distant, with kinematic distances between 1 and 8.5 kpc.  The IRDCs appear to be a very interesting population of molecular clouds. On theoretical grounds, we expect that molecular clouds on the verge of  star formation will contain extremely cold and dense condensations, which will show up as IRDCs if they are favourably located in front of a region of extended mid-IR emission.  The selection criteria for the MSX IRDCs \\citep{car98} ensures that  they do not yet contain high-mass main-sequence stars and HII regions.  Trying to relate the physical properties of this selection of dust clouds to those of better-studied nearby molecular clouds, such as those in the Orion-Taurus region, is problematic.  Most of  the best studied clouds lie in the outer Galaxy where there is insufficient  background emission for IRDCs to be identifiable.  By contrast, the large distances of many MSX IRDCs make them difficult to detect at visible  and near-infrared wavelengths, and difficult to pick out from the surrounding  clutter of lower density clouds in molecular-line surveys.  A notable exception is the cloud G79.3+0.3, located in the Cygnus Rift  at an estimated distance of 800~pc \\citep{mil37, ikh61, wen91}, close enough to the Sun that it can be seen at visible wavelengths \\citep{red00b}. Figure~\\ref{g79} and Figure~\\ref{8mic} show our JCMT SCUBA 850~\\micron\\ image of the G79.3+0.3 IRDC and the MSX 8~\\micron\\ image of the same region, respectively. Unlike the other IRDCs we have observed, this cloud exhibits a variety of indicators of star formation.  Most prominently, the HII region DR15 (G79.307+0.277) lies behind and slightly south of the IRDC, between P4 and P5 in the figure.  It is not immediately obvious whether the IRDC is part of a larger complex including DR15 or is an unrelated foreground cloud. In  visible and near-IR images of the region \\citep{red00b}, there is an association of stars (to the west of DR15) whose northern boundary is defined by the southern edge of the IRDC.  If DR15 is related to this  association, it suggests that both DR15 and the association lie beyond the IRDC.  In the 8~\\micron\\ MSX image (Figure~\\ref{8mic}), the dark filament of  the IRDC (whose dust content is traced by the SCUBA 850~\\micron\\ emission in Figure~\\ref{g79}) is  interrupted between P1 and P3 by a bright patch of warm dust emission which is suggestive of deeply embedded hot stars. Because of this, the eastern and western parts of the IRDC were initially assigned different designations in the MSX catalog (G79.34+0.33 and G79.27+0.38, labeled G79.34 and G79.27 in Figure~\\ref{g79}), although the figure clearly shows them to comprise a single, connected filament. A Herbig-Haro jet that appears to be driven by a YSO at RA(2000) = 20~31~45.5, Dec(2000) = +40~18~44 (see Figure~\\ref{8mic}) has been discovered in front of the IRDC at visible wavelengths \\citep{red00b}.  In this paper we focus our attention on the star-forming activity in the vicinity of  its most prominent condensation, G79.3+0.3~P1. This region was selected for closer study because comparison of the SCUBA and MSX images discussed below showed the presence of a faint, point-like emission source MSX5C G079.3398+00.3415 in the MSX image that was nearly coincident with the brightest compact source of emission at 850 and 450~\\micron . This suggested the presence of a deeply embedded star interacting with the cold dust cloud, but the relatively coarse resolution of both the MSX and SCUBA images hampered a more detailed interpretation. To  address these issues, we combined  JCMT observations at 450~\\micron , BIMA interferometry at 3~mm,  and Gemini North imaging at 10.75~\\micron , together with archival data from the MSX satellite and the 2MASS survey. ",
        "conclusions": "Our Gemini N-band observations  demonstrate that the MSX 8~\\micron\\ source MSX5C G079.3398+00.3415 associated with G79.3+0.3~P1 is dominated by a single compact object that we identify with the star 2MASSI~2032220+402017.  This object is a luminous YSO with a strong IR excess that will likely become an early-B star when it reaches the main sequence.  Two nearby, fainter YSOs were discovered in our N-band images, and a cluster of faint, heavily reddened stars are visible in the 2MASS K-band images of the region.  This concentration of stars and YSOs indicates that the IRDC is actively forming low-mass stars, with the 2MASS star as the most massive star that has formed from it to date.  There are a variety of unusual features in the region around G79.3+0.3~P1 that suggest   2MASSI~2032220+402017 is interacting with the foreground IRDC. Most directly, the BIMA HCO$^+(1-0)$ line has a strongly enhanced blue wing in the velocity range $-1.5$ to 0~km/s, as well as a slight enhancement on the red wing, in a region immediately to the southeast of the 2MASS star.  The emission in the blue wing shows considerable structure on small angular scales.  The unusual nature of this structure is highlighted by the fact that the rest of the cloud shows little emission in the blue wing and very modest small-scale structure in the velocity range $[-0.25, +2.25]$~km/s that traces most of the mass of the cloud.  The presence of small-scale structure close to the 2MASS star is confirmed by CS$~(5-4)$ observations with even higher angular resolution.  The simplest interpretation of these data is that the 2MASS star is a Herbig Ae/Be star that is forming on the far side of the IRDC.  This star must have a warm disk that is the source of the N-band emission, but the disk cannot produce the observed reddening at K.  Hence the star must lie behind a lot of extinction from the IRDC.  The strong wind that would be expected  from such a star is probably impacting the back of the IRDC, exciting the blue wing of the HCO$^+$ line.  Although it appears to be massive, the 2MASS star is still too young to have disrupted the IRDC.  We speculate that in another million years or so this region may resemble a smaller version of DR15 to the south, appearing as an HII region partially obscured by foreground dust left over from the IRDC.  It is, at the very least, a striking coincidence that the two  condensations of dust and gas that comprise the G79.3+0.3~P1 SCUBA source (the A and B components in the BIMA data) lie immediately adjacent in the plane of the sky to the most massive YSO in the cloud, and are apparently being impacted by the same wind that is exciting the blue wing in the HCO$^+(1-0)$ line in the rest of the IRDC.  The combined masses of the two components are comparable to the mass of the Herbig Ae/Be star.  These two objects are the obvious candidates for the next generation of stars to form in the vicinity of the 2MASS star. Further observations may reveal whether the wind from the 2MASS star has played a role in forming the A and B components, or is just now impacting these two pre-existing clumps, possibly driving them into collapse."
    },
    "0212/astro-ph0212458_arXiv.txt": {
        "abstract": "Motivated by the recent direct detection of cosmological gas infall, we develop an analytical model for calculating the mean density profile around an initial overdensity that later forms a dark matter halo. We account for the problem of peaks within peaks; when considering a halo of a given mass we ensure that this halo is not a part of a larger virialized halo. For halos that represent high-sigma fluctuations we recover the usual result that such halos preferentially lie within highly overdense regions; in the limit of very low-sigma fluctuations, on the other hand, we show that halos tend to lie within voids. Combined with spherical collapse, our results yield the typical density and velocity profiles of the gas and dark matter that surround virialized halos. ",
        "introduction": "\\label{sec:Intro}  Numerical simulations of hierarchical halo formation indicate a roughly universal spherically-averaged density profile within virialized dark matter halos \\citep{NFW}. Such profiles are generated by complicated processes that involve violent relaxation and multiple shell crossing; the great interest in them results from the possibility of a direct comparison to the profiles inferred from galactic rotation curves \\citep[e.g.,][]{sb00,dmr01,vs01} and from lensing in clusters \\citep[e.g.,][]{ste02,dhs02}. However, the analogous question of what is the typical density profile outside the virial radius has received relatively little attention, because of the difficulty of directly observing density profiles even out to the halo virial radius; this difficulty results from the rapid decrease in the gas and dark matter densities with distance from the halo center. However, the theoretical description is simpler for infall outside the virialization radius, since in cold dark matter (hereafter CDM) models there is no shell crossing until the final non-linear collapse.  The gas around halos can only be detected using physical effects that are very sensitive to low-density gas. One such phenomenon is resonant \\Lya absorption; \\citet{GP} first noted that even an H I neutral fraction $\\sim 10^{-5}$ suffices to make gas at the cosmic mean density strongly absorb photons at the \\Lya resonance. Thus, the pattern of gas infall around halos may affect observable properties of the \\Lya absorption in these regions. At moderate redshifts, these regions near halos play a prominent role in measurements of the proximity effect of quasars \\citep[e.g.,][]{prox1, prox2}. In these measurements, the effect of the high quasar intensity on the statistics of \\Lya absorption is analyzed and used to infer the mean ionizing background, since the quasar intensity dominates only as long as it is significantly stronger than the background. A different application may be possible at very high redshifts, corresponding to the time before the universe was fully reionized; the presence of the fully neutral IGM could be detected through absorption of a bright source by the \\Lya damping wing of the IGM \\citep{jordi98}. However, every bright source creates a surrounding H II region which complicated the interpretation \\citep{CH00,MR00}. The properties of the gas in the regions that lie just outside the host halo of the source galaxy or quasar should play a major role in these various applications of the observed absorption.  The various calculations of absorption cited above neglected the enhanced densities and the negative (i.e., infall) velocities expected for the gas surrounding a massive halo at high redshift. This infall, however, may be an essential component of any correct interpretation of observations; \\citet{le95} applied it to the proximity effect and showed that neglecting infall can lead to an overestimate of the ionizing background flux by up to a factor of three. Applying the same infall profile, \\citet{nature} have recently revealed the signature of infalling gas around high redshift quasars. The spectral pattern due to \\Lya absorption by the infalling gas can be used to estimate the quasar's ionizing intensity and the total mass of its host halo. While the spectra of individual objects are marked by the large density fluctuations expected in the intergalactic medium (hereafter IGM), if this pattern is averaged over many objects then the resulting mean profile could be compared directly with an accurate theoretical prediction for the density and velocity profile of the infalling gas.  \\citet{le95} defined a halo in the same way as halos were defined in the classic model of \\citet{PS}. In this approach, halos on a given scale are defined as forming when the initial density averaged on this scale is higher than a threshold that depends on the collapse redshift and is fixed using spherical collapse. The halo mass function derived from this assumption, when multiplied by an ad-hoc correction factor of two, provides a reasonable match to the results of numerical simulations \\citep[e.g.,][]{katz93}.  The same mass function was later derived by \\citet{bc91} using a self-consistent approach, with no need for external factors. In their derivation, the factor of two has a more satisfactory origin, namely the so-called ``cloud-in-cloud'' problem: if the average initial density on a certain scale is above the collapse threshold, this region may be contained within a larger region that is itself also denser than the collapse threshold; in this case the original region should be counted as belonging to the halo with mass corresponding to the larger collapsed region. \\citet{bc91} and \\citet{PS} both give the same halo mass function and thus appear to be at least statistically valid given the agreement with the mass function seen in numerical simulations. However, the \\citet{bc91} model is more satisfactory since it accounts for the cloud-in-cloud problem and, more importantly, allows for calculations of other halo properties in addition to simple mass functions, all in an approach that is more self-consistent.  The \\citet{bc91} approach (also known as extended Press-Schechter) has been used to calculate halo merger histories \\citep{lc93}, halo mass functions in models with warm dark matter \\citep{wdm}, halo mass functions in a model with ellipsoidal collapse \\citep{ellipse}, and the nonlinear biasing of halos \\citep{porciani,evan}. In this paper, we present a new application of this approach in which we obtain the expected profile of infalling matter around virialized halos. The rest of this paper is organized as follows. In \\S~2.1 we establish our notation and review the \\citet{bc91} derivation of the \\citet{PS} halo mass function. In \\S~2.2 we review the infall profile used by \\citet{le95} and also derive an alternative model based on similar assumptions but carried out in Fourier space. In \\S~2.3 we derive our model based on the extended Press-Schechter formalism. In \\S~2.4 we then discuss the modified picture of halo virialization once infall is included. In \\S~3 we illustrate the predictions of our model for the initial and final density profiles surrounding halos. Finally in \\S~4 we give our conclusions. ",
        "conclusions": "We have developed a model for calculating the initial density profile around overdensities that later collapse to form virialized halos. This is the first such model to account for the possibility that an overdensity on a given scale may be contained inside a large overdensity on an even bigger scale. As a result, we have found that the mean expected profile [eqs.~(\\ref{eq:ePSk}) and (\\ref{eq:final})] depends on both the mass of the halo and its formation redshift. Starting from this mean initial profile, we have used spherical collapse to obtain the final density profile surrounding the virialized halo.  In reality, there will be some variation in the initial profiles, with each leading to a different final profile. We have estimated the scatter within our spherical model, and found that a sample of around a few dozen halos is required in order to obtain an accurate mean profile by averaging. Halo samples derived from numerical simulations or from observations can be used to test our predictions for the mean and for the scatter of the density and velocity profiles of infalling matter around virialized halos.  Once infall is included, the standard derivation of the mean density enclosed within the virial radius fails. Redefining the virial radius as the radius enclosing a density higher than the cosmic mean by the standard value of $18 \\pi^2$, we have found that the initial overdensity needed for a halo to form at some final redshift is lower than it would have to be for an isolated tophat perturbation [compare eqs.~(\\ref{eq:delc}) and (\\ref{eq:delv})]. Therefore, accounting for infall increases the predicted abundance of rare halos for a given initial power spectrum, but the increase is too small to fully explain the discrepancy between the Press-Schechter prediction for the number of high-mass halos and the number seen in numerical simulations.  We have confirmed that rare halos at a given redshift are surrounded by large, overdense regions of infall, but we have found that the more numerous halos that correspond to low-sigma peaks are surrounded by small infall regions that lie within large voids. However, the infalling gas just beyond the cosmological accretion shock is in general much denser than the cosmic mean; a density larger by a factor of at least 10 occurs around the most massive halos at $z \\ga 2$. Possible applications of our results include the Lyman-$\\alpha$ absorption profiles of high-redshift objects and the proximity effect of quasars at all redshifts."
    },
    "0212/astro-ph0212202_arXiv.txt": {
        "abstract": "AGILE is a small gamma-ray astronomy satellite, with good spatial resolution, excellent timing capabilities and an unprecedented large field of view ($\\sim$1/5 of the sky). It will be the next mission dedicated to high energy astrophysics in the range 30 MeV-50 GeV, and will be launched in 2005. Pulsars are a major topic of investigation of AGILE and, besides studying the small sample of known objects, AGILE will offer the first possibility of detecting several young and energetic radio pulsars that have been discovered since the end of the CGRO mission. We provide an estimate of the expected number of detectable gamma-ray pulsars and present AGILE capabilities for timing analysis with small counting statistics, based on the analysis of data from simulations, from the EGRET archive, and from radio pulsar catalogs. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/astro-ph0212034_arXiv.txt": {
        "abstract": "Absorption and Reprocessing of Gamma-ray burst radiation in the environment of cosmological GRBs can be used as a powerful probe of the elusive nature of their progenitors. In particular, transient X-ray emission line and absorption features in the prompt and early afterglows of GRBs are sensitive to details of the location and density structure of the reprocessing and/or absorbing material. To date, there have been only rather few detections of such features, and the significance is marginal in most individual cases. However, transient X-ray emission lines in GRB afterglows have now been found by four different X-ray satellites, which may justify a more detailed theoretical investigation of their origin. In this paper, I will first present a brief review of the status of observations of transient X-ray emission line and absorption features. I will then discuss general physics constraints which those results impose on isotropy, homogeneity, and location of the reprocessing material with respect to the GRB source, and review the various currently discussed, specific models of GRBs and their environments in which the required conditions could arise. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/hep-th0212290_arXiv.txt": {
        "abstract": "Recent cosmological observations  suggest the existence of a positive \\cc\\ $\\Lambda$ with the magnitude $\\Lambda(G\\hbar/c^3) \\approx 10^{-123}$. This review discusses several aspects of the \\cc\\ both from the cosmological  (sections \\ref{intro}--\\ref{cmbrani}) and field theoretical  (sections \\ref{interpretcc}--\\ref{ccstring})  perspectives. After a brief  introduction to the key issues related to cosmological constant and a  historical overview, a summary of the kinematics and dynamics of the standard Friedmann model of the universe is provided. The observational evidence for cosmological constant,  especially from the supernova results, and the  constraints from the age of the universe, structure formation,   Cosmic Microwave Background Radiation (CMBR) anisotropies and a few others  are described in detail, followed by a discussion of  the theoretical models (quintessence, tachyonic scalar field, ...)    from different perspectives. The latter  part  of the review (sections \\ref{interpretcc}--\\ref{ccstring}) concentrates on more conceptual and fundamental aspects of the \\cc\\  like some alternative interpretations of the \\cc, relaxation mechanisms  to reduce the cosmological constant to the currently observed value, the geometrical structure of the de Sitter spacetime, thermodynamics of the de Sitter universe and the role of string theory in the \\cc\\ problem. ",
        "introduction": "This review discusses several aspects of the \\cc\\ both from the cosmological and field theoretical perspectives with the emphasis  on conceptual and fundamental issues rather than on observational details. The plan of the review is as follows: This section introduces the key issues related to cosmological constant and provides a brief historical overview. (For previous reviews of this subject, from cosmological point of view, see \\cite{jpbr,carolcomb1,carolcomb3,carolcomb2}.) Section \\ref{framework} summarizes the kinematics and dynamics of the standard Friedmann model of the universe paying special attention to features involving the cosmological constant. Section \\ref{evidencecc} reviews the observational evidence for cosmological constant, especially the supernova results, constraints from the age of the universe and a few others. We next study models with evolving cosmological `constant' from different perspectives. (In this review, we shall use the term \\cc\\ in a generalized sense including the scenarios in which cosmological ``constant'' is actually varying in time.)  A phenomenological parameterization is introduced in section  \\ref{paraeqn} to compare theory with observation and is followed up with explicit models involving scalar fields in section \\ref{theorydark}. The emphasis is on quintessence and tachyonic scalar field models and the cosmic degeneracies introduced by them. Section \\ref{sfinuniv} discusses cosmological constant and dark energy in the context of models for structure formation and section \\ref{cmbrani} describes the constraints arising from CMBR anisotropies.  The latter part  of the review concentrates on more conceptual and fundamental aspects of the \\cc. ( For previous reviews of this subject, from  a theoretical physics perspective, see \\cite{swlambda,weincomb1,weincomb2}.) Section \\ref{interpretcc} provides some alternative interpretations of the \\cc\\ which might have a bearing on the possible solution to the problem. Several relaxation mechanisms have been suggested in the literature to reduce the cosmological constant to the currently observed value and some of these attempts are described in section \\ref{relaxcc}. Section \\ref{desittergeom} gives a brief description of the geometrical structure of the de Sitter spacetime and the thermodynamics of the de Sitter universe is taken up in section \\ref{horizons}. The relation between horizons, temperature and entropy are presented at one go in this section and the last section deals with the role of string theory in the \\cc\\ problem.  \\subsection{The many faces of the cosmological constant}\\label{facescc}  Einstein's equations, which determine the dynamics of the spacetime, can be derived from the action (see, eg. \\cite{LL2}): \\begin{equation} A = \\frac{1}{16\\pi G} \\int R \\sqrt{-g} d^4 x + \\int L_{\\rm matter}(\\phi, \\partial \\phi) \\sqrt{-g} d^4x \\label{one} \\end{equation} where $L_{\\rm matter}$ is the Lagrangian for matter depending on some dynamical variables generically denoted as $\\phi$. (We are using units with $c=1$.) The variation of this action with respect to $\\phi$ will lead to the equation of motion for matter $ (\\delta L_{\\rm matter}/\\delta \\phi) =0 $,  in a given background geometry, while the variation of the action with respect to the metric tensor $g_{ik}$ leads to the Einstein's equation \\begin{equation} R_{ik} - \\frac{1}{2} g_{ik} R = 16 \\pi G \\frac{\\delta L_{\\rm matter}}{\\delta g^{ik}} \\equiv 8 \\pi G T_{ik} \\label{three} \\end{equation} where the last equation defines the energy momentum tensor of matter to be $T_{ik} \\equiv 2(\\delta L_{\\rm matter}/\\delta g^{ik})$.  Let us now consider a new matter action $L'_{\\rm matter}=L_{\\rm matter} -(\\Lambda/8\\pi G)$ where $\\Lambda$ is a real constant. Equation of motion for the matter $ (\\delta L_{\\rm matter}/\\delta \\phi) =0 $,  does not change under this transformation since $\\Lambda $ is a constant; but the action now picks up an extra term proportional to $\\Lambda$ \\begin{eqnarray} A  &=&  \\frac{1}{16\\pi G} \\int R \\sqrt{-g} d^4 x + \\int \\left(L_{\\rm matter} -\\frac{\\Lambda}{8\\pi G}\\right) \\sqrt{-g} d^4x\\nonumber \\\\ &=&  \\frac{1}{16\\pi G} \\int (R-2\\Lambda) \\sqrt{-g} d^4 x+ \\int L_{\\rm matter} \\sqrt{-g} d^4x \\label{four} \\end{eqnarray} and equation (\\ref{three}) gets modified. This  innocuous looking addition of a constant to the  matter Lagrangian leads to one of the most fundamental and fascinating problems of theoretical physics. The nature of this problem and its theoretical backdrop acquires  different shades of meaning depending which of the two forms of equations in  (\\ref{four}) is used.  The first interpretation, based on the first line of equation (\\ref{four}), treats $\\Lambda $  as the shift in the matter Lagrangian which, in turn, will lead to a shift in the matter Hamiltonian. This could be thought of as a shift in the zero point energy of the matter system. Such a constant shift in the energy does not affect the dynamics of matter while gravity --- which couples to the total energy of the system --- picks up an extra contribution in the form of a new term $Q_{ik}$ in the energy-momentum tensor, leading to: \\begin{equation} R^i_k - \\frac{1}{2} \\delta^i_k R = 8\\pi G (T^i_k + Q^i_k); \\qquad Q^i_k \\equiv \\frac{\\Lambda}{8 \\pi G} \\delta^i_k \\equiv \\rho_\\Lambda \\delta^i_k \\label{six} \\end{equation}  The second line in equation (\\ref{four}) can be interpreted as gravitational field, described by the Lagrangian of the form $L_{\\rm grav} \\propto (1/G) (R-2\\Lambda)$, interacting with matter described by the Lagrangian $L_{\\rm matter}$. In this interpretation, gravity is described by two constants, the Newton's constant $G$ and the {\\it \\cc}\\  $\\Lambda$. It is then natural to modify the {\\it left hand side} of Einstein's equations and write (\\ref{six}) as: \\begin{equation} R^i_k - \\frac{1}{2} \\delta^i_k R - \\delta^i_k \\Lambda = 8\\pi G T^i_k \\label{five} \\end{equation} In this interpretation, the spacetime is treated as curved even in the absence of matter ($T_{ik} =0$) since the equation $R_{ik} - (1/2) g_{ik}R -\\Lambda g_{ik} =0$ does not admit flat spacetime as a solution. (This situation is rather unusual and is related to the fact that symmetries of the theory with and without a \\cc\\ are drastically different; the original symmetry of general covariance cannot be naturally broken in such a way as to preserve the sub group of spacetime translations.)  In fact, it is possible to consider a situation in which both effects can occur. If the gravitational interaction is actually described by the Lagrangian of the form $(R-2\\Lambda)$, then there is an intrinsic cosmological constant in nature just as there is a Newtonian gravitational constant in nature. If the matter Lagrangian contains energy densities which change due to dynamics, then $L_{\\rm matter}$ can pick up constant shifts during dynamical evolution. For example, consider a scalar field with the Lagrangian $L_{\\rm matter} =(1/2) \\partial_i\\phi \\partial^i\\phi - V(\\phi)$ which has the energy momentum tensor \\begin{equation} T^a_b = \\partial^a\\phi \\partial_b \\phi - \\delta^a_b \\left( \\frac{1}{2} \\partial^i\\phi\\partial_i\\phi - V(\\phi)\\right) \\label{scalartab} \\end{equation} For field configurations which are constant [occurring, for example, at the minima of the potential $V(\\phi)$], this contributes an energy momentum tensor $T^a_b = \\delta^a_b V(\\phi_{\\rm min})$ which has exactly the same form as a \\cc. As far as gravity is concerned, it is the combination of these two effects --- {\\it of very different nature} --- which is relevant and the source will be $T_{ab}^{\\rm eff} = [V(\\phi_{\\rm min}) + (\\Lambda/8\\pi G)]g_{ab}$, corresponding to an effective gravitational constant \\begin{equation} \\Lambda_{\\rm eff} = \\Lambda + 8\\pi G V(\\phi_{\\rm min}) \\label{nine} \\end{equation} If $\\phi_{\\rm min}$ and hence $V(\\phi_{\\rm min})$ changes during dynamical evolution, the value of $\\Lambda_{\\rm eff}$ can also change in course of time. More generally, any field configuration which is varying slowly in time will lead to a slowly varying $\\Lambda_{\\rm eff}$.   The extra term $Q_{ik}$ in Einstein's equation  behaves in a manner which is very peculiar compared to the energy momentum tensor of normal matter. The term $Q^i_k = \\rho_\\Lambda \\delta^i_k$ is in the form of the energy momentum tensor of an ideal fluid with energy density $\\rho_\\Lambda$ and pressure $P_\\Lambda = - \\rho_\\Lambda$; obviously, either the pressure or the energy density of this ``fluid'' must be negative, which is unlike conventional laboratory systems. (See, however, reference \\cite{volovik}.)  Such an equation of state, $\\rho = -P $ also has another important implication in general relativity. The spatial part  ${\\bf g}$  of the geodesic acceleration (which measures the relative acceleration of two geodesics in the spacetime) satisfies the following exact equation in general relativity (see e.g., page 332 of \\cite{probbook}): \\begin{equation} \\nabla \\cdot {\\bf g} = - 4\\pi G (\\rho + 3P) \\label{nextnine} \\end{equation} showing that the source of geodesic  acceleration is $(\\rho + 3P)$ and not $\\rho$. As long as $(\\rho + 3P) > 0$, gravity remains attractive while $(\\rho + 3P) <0$ can lead to repulsive gravitational effects. Since the \\cc\\ has $(\\rho_\\Lambda + 3P_\\Lambda) = -2\\rho_\\Lambda$, a positive \\cc\\ (with $\\Lambda >0$) can lead to {\\it repulsive} gravity.  For example, if the energy density of normal, non-relativistic  matter with zero pressure is $\\rho_{\\rm NR}$, then  equation  (\\ref{nextnine}) shows that the geodesics will accelerate away from each other due to the repulsion of \\cc\\ when $\\rho_{\\rm NR} < 2\\rho_\\Lambda$. A related feature, which makes the above conclusion practically relevant is the fact that, in an expanding   universe, $\\rho_\\Lambda$ remains constant while $\\rho_{\\rm NR}$ decreases. (More formally, the equation of motion, $d(\\rho_\\Lambda V) = - P_\\Lambda dV$ for the cosmological constant, treated as an ideal fluid, is identically satisfied with constant $\\rho_\\Lambda, P_\\Lambda$.) Therefore, $\\rho_\\Lambda $ will eventually dominate over $\\rho_{\\rm NR}$ if the universe expands sufficiently. Since $|\\Lambda|^{1/2}$ has the dimensions of inverse length, it will set the scale for the universe when cosmological constant dominates.  It follows that the most stringent bounds on $\\Lambda$ will  arise from cosmology when the expansion of the universe has diluted the matter energy density sufficiently. The rate of expansion of the universe today is usually expressed in terms of the Hubble constant: $H_0 = 100 h $ km s$^{-1}$ Mpc$^{-1}$ where 1 Mpc  $ \\approx 3\\times 10^{24}$ cm and  $h$ is a dimensionless parameter in the range $0.62\\lesssim h \\lesssim 0.82$ (see section \\ref{ageuniv}). From $ H_0$ we can form the time scale $t_{\\rm univ}\\equiv H_0^{-1}\\approx 10^{10} h^{-1}$ yr and the length scale $  cH_0^{-1} \\approx 3000 h^{-1}$ Mpc; $t_{\\rm univ}$  characterizes the evolutionary time scale of the universe and $H_0^{-1}$ is of the order of the largest length scales currently accessible in cosmological observations. From the observation that the universe is at least of the size $H_0^{-1}$, we can set a bound on $\\Lambda $ to be $|\\Lambda| < 10^{-56}$ cm$^{-2}$. This stringent bound  leads to several issues which have been debated for decades without satisfactory solutions. \\begin{itemize} \\item In classical general relativity, based on the constants $G, c $ and $\\Lambda$,  it is not possible to construct any dimensionless combination from these constants. Nevertheless, it is clear that $\\Lambda$ is extremely tiny compared to any other physical scale in the universe,  suggesting that $\\Lambda$ is probably zero. We, however, do not know of any symmetry mechanism or invariance principle which requires $\\Lambda$ to vanish. Supersymmetry does require the vanishing of the ground state energy; however, supersymmetry is so badly broken in nature that this is not of any practical use \\cite{supersym1,supersym2}.  \\item We mentioned above that observations actually constrain $\\Lambda_{\\rm eff}$ in equation (\\ref{nine}), rather than $\\Lambda$. This requires   $\\Lambda$ and $V(\\phi_{\\rm min})$ to be fine tuned to an enormous accuracy for the bound, $|\\Lambda_{\\rm eff}| < 10^{-56} \\ {\\rm cm}^{-2}$,  to be satisfied. This becomes more mysterious when we realize that $V(\\phi_{\\rm min})$ itself could change by several orders of magnitude during the evolution of the universe.  \\item When quantum fields  in a given curved spacetime are considered (even without invoking any quantum gravitational effects) one introduces the Planck constant, $\\hbar$, in the description of the physical system. It is then possible to form the dimensionless combination $\\Lambda (G\\hbar/c^3) \\equiv \\Lambda L_P^2$. (This equation also defines the quantity $L_P^2$; throughout the review we use the symbol `$\\equiv$' to define variables.) The bound on $\\Lambda$ translates into the condition $\\Lambda L_P^2 \\lesssim 10^{-123}$. As has been mentioned several times in literature, this will require enormous fine tuning.  \\item All the above questions could have been satisfactorily answered if we take $\\Lambda_{\\rm eff}$ to be zero and assume that the correct theory of quantum gravity will provide an explanation for the vanishing of \\cc. Such a view was held by several people (including the  author) until very recently. Current cosmological observations however suggests that $\\Lambda_{\\rm eff}$ is actually non zero and $\\Lambda_{\\rm eff} L_P^2$ is indeed of order ${\\mathcal O}(10^{-123})$. In some sense, this is the cosmologist's worst nightmare come true. {\\it  If the observations are correct, then $\\Lambda_{\\rm eff}$ is non zero, very tiny and its value is extremely fine tuned for no good reason}. This is a concrete statement of the first of the two `\\cc\\ problems'.  \\item The bound on $\\Lambda L_P^2$ arises from the expansion rate of the universe or --- equivalently --- from the energy density which is present in the universe today. The observations require the energy density of normal, non relativistic matter to be of the same order of magnitude as the energy density contributed by the cosmological constant. But in the past, when the universe was smaller, the energy density of normal matter would have been higher while the energy density of \\cc\\ does not change.  Hence we need to adjust the energy densities of normal matter and \\cc\\ in the early epoch very carefully so that $\\rho_\\Lambda\\gtrsim \\rho_{\\rm NR}$ around the current epoch. If this had happened very early in the evolution of the universe, then the repulsive nature of a positive cosmological constant would have initiated a rapid expansion of the universe,  preventing the formation of galaxies, stars etc. If the epoch of $\\rho_\\Lambda \\approx \\rho_{\\rm NR}$ occurs much later  in the future, then the current observations would not have revealed the presence of non zero \\cc. This raises the second of the two \\cc\\ problems: Why is it that $(\\rho_\\Lambda/ \\rho_{\\rm NR}) = \\mathcal{O} (1)$ at the {\\it current} phase of the universe ?  \\item The sign of $\\Lambda$ determines the nature of solutions to Einstein's equations as well as the sign of $(\\rho_\\Lambda + 3P_\\Lambda)$. Hence the spacetime geometry with $\\Lambda L_P^2 = 10^{-123}$ is very different from the one with $\\Lambda L_P^2 = - 10^{-123}$. Any theoretical principle which explains the near zero value of $\\Lambda L_P^2$ must also explain why the observed value of $\\Lambda$ is positive.  \\end{itemize}  At present we have no clue as to what the above questions mean and how they need to be addressed. This review summarizes different attempts to understand the above questions from various perspectives.     \\subsection{A brief history of \\cc} \\label{history}  Originally, Einstein introduced the \\cc\\ $\\Lambda$ in the field equation for gravity (as in equation (\\ref{five})) with the motivation that it allows for a finite, closed, static universe in which the energy density of matter determines the geometry. The spatial sections of such a universe are closed 3-spheres with radius $l = (8\\pi G \\rho_{\\rm NR})^{-1/2} = \\Lambda^{-1/2}$ where $\\rho_{\\rm NR}$ is the energy density of pressureless matter (see section \\ref{geometry}) Einstein had hoped that   normal matter is {\\it needed} to curve the geometry;   a demand, which ---  to him ---  was closely related to the Mach's principle. This hope, however, was soon shattered when de Sitter produced a solution to Einstein's equations with \\cc\\ containing no matter \\cite{desitter}. However,  in spite of two fundamental papers by Friedmann and one by Lemaitre \\cite{F1,F2}, most workers did not catch on with the idea of an expanding universe. In fact, Einstein originally thought Friedmann's work was in error but later published a retraction of his comment; similarly, in the Solvay meeting in 1927, Einstein was arguing against the solutions describing expanding universe. Nevertheless, the Einstein archives do contain a postcard from Einstein to Weyl in 1923 in which he says: ``If there is no quasi-static world, then away with the cosmological term''. The early history following de Sitter's discovery (see, for example, \\cite{north}) is clearly somewhat confused,  to say the least.  It appears that the community accepted the concept of an expanding universe largely due to the work of Lemaitre. By 1931, Einstein himself had rejected the cosmological term as superflous and unjustified (see reference \\cite{einstein}, which is a single authored paper; this paper has been mis-cited in literature often, eventually converting part of the journal name ``preuss'' to a co-author ``Preuss, S. B''!; see \\cite{history}). There is no direct record that Einstein ever called \\cc\\ his biggest blunder. It is possible that this often repeated ``quote'' arises from Gamow's recollection \\cite{gamow}: ``When I was discussing cosmological problems with Einstein, he remarked that the introduction of the cosmological term was the biggest blunder he ever made in his life.'' By 1950's  the view was decidedly against $\\Lambda$ and the authors of several  classic texts ( like Landau and Liftshitz\\cite{LL2}, Pauli \\cite{pauli} and Einstein \\cite{meaning}) argued against the \\cc.  In later years, \\cc\\ had a chequered history and was often accepted or rejected for  wrong or insufficient reasons. For example, the original value of the Hubble constant was nearly an order of magnitude higher \\cite{hubble29} than the currently accepted value thereby reducing the age of the universe by a similar factor. At this stage, as well as on several later occasions (eg., \\cite{sandage61,gunntinsley}), cosmologists have invoked \\cc\\ to reconcile the age of the universe with observations (see section \\ref{ageuniv}). Similar attempts have been made in the past when it was felt that counts of quasars peak at a given phase in the expansion of the universe \\cite{petro,shklov,kardash}. These reasons, for the introduction of something as fundamental as \\cc, seem inadequate at present.  However, these attempts clearly showed that sensible cosmology can only be obtained if the energy density contributed by \\cc\\ is comparable to the energy density of matter at the present epoch. This remarkable property was probably noticed first by Bondi \\cite{bondi} and has been discussed by McCrea \\cite{mccrea}. It is mentioned in \\cite{jpbr} that such coincidences were discussed in Dicke's gravity research group in the sixties; it is almost certain that this must have been noticed by several other workers in the subject.  The first cosmological model to make central use of the \\cc was the steady state model \\cite{steady1,steady2,steady3}. It made use of the fact that a universe with a cosmological constant has a time translational invariance in a particular coordinate system. The model also used a scalar field with negative energy field to continuously create matter while maintaining energy conservation. While modern approaches to cosmology invokes negative energies or pressure without hesitation, steady state cosmology was discarded by most workers after the discovery of CMBR.  The discussion so far has been purely classical. The introduction of quantum theory adds a new dimension to this problem. Much of the early work \\cite{earlywork1a,earlywork1b} as well as the definitive work by Pauli \\cite{earlywork2a,earlywork2b} involved evaluating the sum of the zero point energies of a quantum field (with some cut-off) in order to estimate the vacuum contribution to the \\cc. Such an argument, however,  is hopelessly naive (inspite of the fact that it is often repeated even today). In fact, Pauli himself was aware of the fact that one must {\\it exclude} the zero point contribution from such a calculation. The first paper to stress this clearly and carry out a {\\it second order} calculation was probably the one by Zeldovich \\cite{zeldo} though the connection between vacuum energy density and cosmological constant had been noted earlier by Gliner \\cite{gliner} and even by Lemaitre \\cite{earlywork3}. Zeldovich assumed that the lowest order zero point energy should be subtracted out in quantum field theory and went on to compute the gravitational force between particles in the vacuum fluctuations. If $E$ is an energy scale of a virtual process corresponding to a length scale $l=\\hbar c/E$, then   $l^{-3}=(E/\\hbar c)^3$ particles  per unit volume of energy $E$ will lead to the gravitational self energy density of the order of \\begin{equation} \\rho_\\Lambda \\approx \\frac{G(E/c^2)^2}{l}l^{-3} = \\frac{GE^6}{c^8\\hbar^4} \\end{equation} This will correspond to $\\Lambda L_P^2\\approx (E/E_P)^6$ where $E_P=(\\hbar c^5/G)^{1/2}\\approx 10^{19}$GeV is the Planck energy. Zeldovich took $E\\approx 1$ GeV (without any clear reason) and obtained a $\\rho_\\Lambda$ which contradicted the observational bound  ``only'' by nine orders of magnitude.  The first serious symmetry principle which had implications for \\cc\\ was supersymmetry and it was realized early on \\cite{supersym1,supersym2} that the contributions to vacuum energy from fermions and bosons will cancel in a supersymmetric theory. This, however, is not of much help since supersymmetry is badly broken in nature at sufficiently high energies (at $E_{\\rm SS} > 10^2$ Gev). In general, one would expect the vacuum energy density to be comparable to the that corresponding to the supersymmetry braking scale, $E_{\\rm SS}$. This will, again, lead to an unacceptably large value for $\\rho_\\Lambda$. In fact the situation is more complex and one has to take into account the coupling of matter sector and gravitation --- which invariably leads to a supergravity theory. The description of \\cc\\ in such models is more complex, though none of the attempts have  provided a clear direction of attack (see e.g, \\cite{swlambda} for a review of early attempts).  The situation becomes more complicated when the quantum field theory admits more than one ground state or even more than one local minima for the potentials. For example, the spontaneous symmetry breaking in the electro-weak theory arises from a potential of the form \\begin{equation} V = V_0 - \\mu^2  \\phi^2 + g \\phi^4 \\qquad (\\mu^2, g>0) \\end{equation} At the minimum, this leads to an energy density $V_{\\rm min} = V_0 - (\\mu^4/4g)$. If we take $V_0 =0$ then $(V_{\\rm min}/g) \\approx -(300\\ {\\rm GeV})^4$. For an estimate, we will assume that the gauge coupling constant $g$ is comparable to the electromagnetic coupling constant: $g ={\\mathcal O}(\\alpha^2)$, where $\\alpha \\equiv (e^2/\\hbar c)$ is the fine structure constant. Then, we get $|V_{\\rm min}| \\sim 10^6\\ {\\rm GeV}^4$ which misses the bound on $\\Lambda$ by a factor of $10^{53}$. It is really   of no help to set $V_{\\rm min}=0$ by hand. At early epochs of the universe, the temperature dependent effective potential  \\cite{thermo1,thermo2} will change minimum to $\\phi =0$ with $V(\\phi) = V_0$. In other words, the ground state energy changes by several orders of magnitude during the electro-weak and other phase transitions.  Another facet is added to the discussion by the currently popular models of quantum gravity based on string theory \\cite{stringtext1,stringtext2}. The currently accepted  paradigm of string theory encompasses several ground states of the same underlying theory (in a manner which is as yet unknown). This could lead to the possibility that the final theory of quantum gravity might allow different ground states for nature and we may need an {\\it extra} prescription to choose the actual  state in which we live in. The  different ground states can also have different values for \\cc\\ and we need to invoke a separate (again, as yet unknown) principle to choose the ground state in which $\\Lambda L_P^2 \\approx 10^{-123}$ (see section \\ref{ccstring}). ",
        "conclusions": ""
    },
    "0212/astro-ph0212278_arXiv.txt": {
        "abstract": "This poster discusses a possible explanation for the relationship between the mass of the central supermassive black hole and the velocity dispersion in the bulge of the host galaxy.  We suppose that the black hole and the dark matter halo are forming simultaneously as matter falls in, and a self-similar system then exists in which the mass and the velocities of the system evolve as power-law functions of time.  This leads naturally to a relationship between the black hole mass and the velocities in the halo which, with a reasonable choice of cosmological parameters, is in good agreement with the observed relationship.  We also confirm this relationship with more robust numerical results. ",
        "introduction": "Supermassive black holes (BHs) are now considered to be a common feature of galaxies which have a bulge.  Furthermore, a number of observational properties of the host galaxy correlate with the BH mass.  Among the strongest of these correlations is the relationship between the BH mass and velocity dispersion within the galactic bulge (Ferrarese \\& Merritt 2000; Gebhardt et al. 2000):  $M_{\\rm BH} \\propto \\sigma^a$, where $a = 4.02 \\pm 0.32$ (Tremaine et al. 2002).  Since the velocity dispersion is measured well outside of the BH ``influence,'' this correlation indicates that an intimate relationship exists between the BH and the dynamical structure of the host galaxy.  We assume that a galaxy forms by the extended collapse of a ``halo'' composed of collisionless matter and that simultaneously the central black hole is growing proportionally to the halo as matter continues to fall in.  This is equivalent to the assumption of multidimensional self-similarity (Carter \\& Henriksen 1991; Henriksen 1997), and MacMillan \\& Henriksen (2002) show that, in this case, the mass $M$ inside any surface in the halo is related to the velocity of the particles by \\begin{equation} \\log{ M}\\propto \\left(\\frac{3\\delta/\\alpha-2}{\\delta/\\alpha-1}\\right) \\log{\\sigma}, \\label{eq:predict} \\end{equation} where $\\sigma$ is the averaged velocity.  Since the system is self-similar, this relation will apply at both the ``bulge'' mass scale as well as the BH mass scale.  The quantities $\\delta$ and $\\alpha$ are scales in space and time, respectively, and are related to the power-law index of the initial density perturbation $\\epsilon$ in spherical infall models of halo growth (Henriksen \\& Widrow 1999): \\begin{equation} \\frac{\\delta}{\\alpha}=\\frac{2}{3}\\left(1+\\frac{1}{\\epsilon}\\right). \\label{eq:relate} \\end{equation}  This power-law index $\\epsilon$ depends on the spectral index $n$ of the primordial power spectrum, $P(k) \\propto k^n$, by $n = 2\\epsilon - 3$.  There is therefore a direct link between the initial  power spectrum and the predicted relationship between the BH mass and the velocity dispersion; for agreement with observations, we must have $n = -2$, so that $\\epsilon = 1/2$ and $\\delta/\\alpha = 2$.  We determine below that the addition of sufficient angular momentum breaks the self-similarity in the central regions.  This creates a ``mass reservoir'' around the BH that does grow in proportion to the galaxy mass, and from which the BH grows more slowly by collisional interactions between clumps of matter.  Provided most of this mass may be absorbed by the BH on a cosmic timescale, the proposed relation should still hold. ",
        "conclusions": "For systems in which there is no angular momentum or deviation from spherical symmetry initially, the numerical work, with both the spherical shell code and the more general n-body code, confirms the predicted black hole growth:  $M_{\\rm BH} \\propto t^4$.  However, for the simulations in which there was some deviation from a simple spherical infall of matter, results were contrary to what was expected.  In particular, the n-body results show a black hole growth approximately proportional to the time.  However, it is apparent from the simulations that the black hole growth in these cases is not dominated by the infalling matter, as it is in the initially spherical case.  Rather, the particles form a ``core'' about the centre, and so the growth is fed in a different manner.  Dimensional arguments suggest that the Schwarzschild radius should scale with time as $R_s \\propto ct$, where $c$ is the speed of light, a constant of the system.  If this is the case, then we get that the black hole mass will grow as $M_{\\rm BH} \\propto R_s \\propto t$, which is approximately what is observed.  Note, however, that this breaks the symmetry of the system, since the mass on a larger scale is predicted to grow as $t^{3\\delta - 2\\alpha}$.  The numerically derived $M_{\\rm BH} - \\sigma$ relationship, which here compared the black hole growth with the average velocity of particles within some radius $r$, only gave the expected relation, $M_{\\rm BH} \\propto \\sigma^4$, when a fixed, rather than comoving radius was considered.  Furthermore, the cases for which the growth rate was much shallower than expected still showed the same trend as the other case; that is, $M_{\\rm BH} \\propto \\sigma^4$ regardless of how the black hole grew with time.  As stated above, however, this result is for a fixed radius.  From the definition of $\\mathbf{X}$ and $T$ in MacMillan \\& Henriksen (2002), we see that $r = \\mathbf{X} t^{\\delta / \\alpha}$.  Fixing $r$ requires $\\mathbf{X}$ to then depend on time.  Recalling that we're considering a density profile that goes as $X^{-\\epsilon}$, and $\\epsilon$ is related to the quantity $\\delta / \\alpha$ by equation (\\ref{eq:relate}), we can write the mass as a function of $X$: \\begin{equation} \\mathcal{M}(X) \\propto X^{3 - 2 \\alpha / \\delta}. \\end{equation} Thus, as a function of the (fixed) radius, we have $M(r) \\propto r^{3 - 2 \\alpha / \\delta}$.  Similarly, we can write the self-similar velocity, assuming it is in its virialized state, as a function of $X$: \\begin{equation} Y(X) \\propto \\sqrt{\\bar{\\Phi}} \\propto X^{1 - \\alpha / \\delta}, \\end{equation} so that the radial dependence of the velocity takes the form $v(r) \\propto r^{1 - \\alpha / \\delta}$.  Combining these two equations for mass and velocity and eliminating $r$ gives us the familiar relation given by equation (\\ref{eq:predict}).  Thus, taking the mass and velocity at a fixed radius $r$ gives the same relation regardless of the time dependence of either."
    },
    "0212/astro-ph0212087_arXiv.txt": {
        "abstract": "We combine Sloan Digital Sky Survey spectra of 22,000 luminous, red, bulge-dominated galaxies to get high $S/N$ average spectra in the rest-frame optical and ultraviolet (2600\\,\\AA\\ to 7000\\,\\AA). The average spectra of these massive, quiescent galaxies are early-type with weak emission lines and with absorption lines indicating an apparent excess of $\\alpha$ elements over solar abundance ratios.  We make average spectra of subsamples selected by luminosity, environment and redshift.  The average spectra are remarkable in their similarity.  What variations do exist in the average spectra as a function of luminosity and environment are found to form a nearly one-parameter family in spectrum space.  We present a high signal-to-noise ratio spectrum of the variation.  We measure the properties of the variation with a modified version of the Lick index system and compare to model spectra from stellar population syntheses.  The variation may be a combination of age and chemical abundance differences, but the conservative conclusion is that the quality of the data considerably exceeds the current state of the models. ",
        "introduction": "Luminous elliptical and bulge-dominated galaxies are the most massive galaxies in the Universe.  These objects show little rotation, have smooth radial profiles, are massive and kinematically hot, have a narrow range of stellar-mass-to-light ratios, show high metallicities, and reside preferentially in the Universe's denser environments \\citep[e.g.,][]{Kor89,Rob94,Ben98}.  These galaxies are also the most intrinsically luminous galaxies in the Universe \\citep[e.g.,][]{Tam79,Bla01a}. That the most optically luminous galaxies are red is remarkable, given that unextincted red stellar populations generally have higher stellar mass than blue populations at constant luminosity; it means that they are considerably more massive than their bluer, fainter spiral neighbors.  A common paradigm for the formation of bulge-domi\\-nated systems is hierarchical merging of smaller, star-forming progenitors \\citep{Bar92,Kau96,Bau96}. In this model, the systems appear old and metal-rich because the merging is more common at high redshift and because the resulting starburst consumes or expels the remaining gas, thereby ending star formation \\citep{Kau98}. This dependence upon cosmological merger rates and feedback suggests that the stellar populations of bulges should display subtle but important dependencies on mass and environment.  In this paper, we study the aggregate stellar populations of luminous, red, bulge-dominated galaxies using averaged spectra from the Sloan Digital Sky Survey \\citep[SDSS;][]{Yor00} We select these galaxies because we know that they dominate the stellar mass density of the Universe \\citep{Fuk98,Hog02} and because they show great regularities in their properties \\citep[e.g.,][]{Fab73,Vis77,Djo87,Dre87,Kor89,Bow92,Rob94, Ber01}. In addition, the great age of their stellar populations ease some, but not all, of the difficulties in understanding superimposed younger populations \\citep{Fab72,OCo76,Gun81,Cha96}.  As the luminosity-weighted average spectrum does measure the total stellar population of an entire sample of galaxies, it is a well-defined mean property that can be compared to theoretical models of galaxy formation \\citep[][e.g.,]{Bal02}. We focus not on the interpretations of the average spectrum itself, but rather on the variations of this average with luminosity, environment, and redshift.  These variations should be able to test predictions that more luminous (or massive) galaxies have higher metallicity and galaxies in clusters have older stellar populations. Previous observational studies \\citep[e.g.,][]{Bow90,Guz91,Ros94,Jam99,Ter99,Pog01} comparing spectra of individual galaxies have found evidence to support these hypotheses. Of course, one would also expect to find that galaxies at high redshift should have younger stellar populations than similar galaxies at low redshift, but even this trend can be altered if massive, bulge-dominated galaxies are still forming at low redshift \\citep{Fra95, van01}.   Because we are interested in comparing the average spectra of samples, it is critical that the samples be selected so that the variation between the samples is as controlled as possible. With SDSS data, we can select uniform samples according to rest-frame photometry and morphological properties across a wide range of environments and redshifts, allowing trends in the data to be attributed to astrophysics rather than selection biases.   The average spectra are of extremely high signal-to-noise ratio, thereby allowing analyses that would be impractical on individual SDSS spectra. Over 100 absorption lines are apparent in the composite. Indeed, the resulting spectra exceed the ability of the best spectral models of stellar populations to interpret, especially given the non-solar ratio of $\\alpha$-element to iron-peak abundances \\citep{Wor92,Gon93,Dav93,Paq94,Wor98,Pro02} apparent in the data.   The outline of this paper is as follows.  In \\S~2, we describe the relevant aspects of SDSS data and introduce the methods for our analysis, including principal component analysis (PCA) and a modification of the Lick absorption line index system.  In \\S~3, we describe our sample selection and environment estimation.    We present a general analysis of the average spectra in \\S~4 and analysis of the variations with luminosity and environment in \\S~5.  Variations with redshift are presented in \\S~6.  In \\S~7, we use the highest redshift sample to determine the average spectrum in the mid-ultraviolet. We conclude in \\S~8.  Except where noted otherwise, we adopt a conventional cosmological world model with $H_0=100\\,h~{\\rm km\\,s^{-1}\\,Mpc^{-1}}$ and $(\\Omega_M,\\Omega_{\\Lambda})=(1/3,2/3)$. ",
        "conclusions": ""
    },
    "0212/astro-ph0212108_arXiv.txt": {
        "abstract": "The origin of large scale cosmological magnetic fields remains a mystery, despite the continuous efforts devoted to that problem. We present a new model of magnetic field generation, based on local charge separation provided by an anisotropic and inhomogeneous radiation pressure. In the cosmological context, the processes we explore take place at the epoch of the reionisation of the Universe. Under simple assumptions, we obtain results (i) in terms of the order of magnitude of the field generated at large scales and (ii) in terms of its  power spectrum. The amplitudes obtained ($B\\sim 8.~10^{-6} ~\\mu{\\rm G}$) are considerably higher than those obtained in usual magnetogenesis models and provide suitable seeds for amplification by adiabatic collapse and/or dynamo during structure formation. % ",
        "introduction": "According to observations, magnetic fields are pervading the whole Universe. In particular, both synchrotron emission and Faraday rotation measurements reveal a magnetic field present on large scales. On galactic scales \\cite{beck}, on galaxy cluster scales \\cite{eo} and even on larger scales \\cite{kronberg}, the observed magnetic field strengths range from $0.1$ micro-Gauss to a few tens of micro-Gauss.  The origin of the magnetic field is, however, still a mystery. Clearly, the fields cannot be generated directly with the observed amplitudes since, in that case, they would strongly affect the formation of structures \\cite{rr}. Therefore, the fields must arise in the form of weak seeds which are subsequently, in some way, strongly amplified after the onset of structure formation. Various mechanisms have been proposed for the generation of such seeds. Basically, they can be sorted in two types.  In cosmology, the mechanisms of the first type operate before matter-radiation decoupling. They take advantage of high energy physics processes such as inflation \\cite{ratra,dimo,maroto} or phase transitions in the early Universe \\cite{sigl,grario}. However, these mechanisms are still unsatisfactory in two respects. First, they do not provide a real prediction since the generated fields can be as weak as $10^{-65}$ Gauss and as strong as $10^{-9}$ Gauss. Moreover, a property shared by fields created before decoupling is that the stronger their amplitude, the smaller the scale on which they are generated (see Battaner \\& Lesch~\\cite{bl}, and references therein). Furthermore, under the assumption of \\textit{frozen magnetic flux}, these fields decay following cosmological expansion as $(1+z)^2$. Therefore, fields of primordial origins (generated before nucleosynthesis, say) are extremely weak on galactic scales and are very strong on much smaller scales. Finally, extremely tight constraints on pre-BBN magnetic fields from gravity wave production have been recently derived \\cite{chiara}, possibly ruling out a majority of the proposed models.  The mechanisms of the second type produce magnetic fields after matter-radiation decoupling. They rely essentially on the battery effect \\cite{battery} which was first introduced to explain stellar magnetic fields and is basically a consequence of charge separation. Several authors explored the same mechanism in various cosmological contexts \\cite{lazarian,lc,khanna}. The intensity of the fields generated is found to be roughly of the order of $10^{-20}$ Gauss. These weak seeds are then amplified on galactic scales by a dynamo sustained by turbulence and differential rotation in protogalaxies. As noticed by Grasso \\& Rubinstein \\cite{GR}, this type of mechanism can hardly be responsible for magnetic fields on galaxy cluster scales. Consequently, we would have to deal with different mechanisms to explain the origin of fields in galaxies and clusters, which exhibit similar properties. Moreover, this mechanism cannot really account for the strong fields observed in objects at high redshifts where galaxies did not have enough time to rotate and amplify the seeds efficiently.  Finally, whatever the problems of each type of mechanisms, the seeds produced are generally small. As such, they require the action of a strong process capable of amplifying the magnetic amplitudes up to the values observed in galaxies.  The dynamo mechanism evoked above has long been considered suitable for that purpose but it is, however, subject to controversy  nowadays. In fact, it has recently been pointed out \\cite{ketal} that dynamo  tends to amplify the small scale magnetic fields in a more efficient way than fields on  large scales. The consequence of this is that small scale fields reach rapidly dynamically important strengths and stop, in back-reaction, the dynamo mechanism itself. This happens long before large scale fields can reach the observed equipartition values.   Significant progress has been achieved very recently and the back-reaction of small scale fields is now better understood. In particular, the importance of magnetic helicity conservation in dynamo quenching has been demonstrated in ideal cases \\cite{fb02,blackman,branden}. However, for their application to real astrophysical systems such as galaxies, these theories require further developments. Small scale fields thus remain a possible threat to dynamo amplification mechanisms.   In this paper, we propose a mechanism which appears to be relatively free of the problems outlined above. Applied to cosmology, it operates after decoupling in the natural context of inhomogeneous reionisation, for a redshift $z$ between 6 and 7 \\cite{becker, go}. Similarly to the battery effect, it relies upon local charge separation provided in our case by radiation pressure. The physics involved is fairly simple and was nevertheless left aside until now, surprisingly enough. Notice however that Subramanian \\textit{et. al.}~\\cite{subra} proposed  a mechanism (which has been explored numerically by Gnedin \\textit{et. al.}~\\cite{gnedin}) for the generation  of magnetic field seeds at reionisation, and obtained magnetic fields of the order of $10^{-20} - 10^{-19}$ Gauss on galactic scales. Their solution for magnetogenesis relies explicitly on the actual battery effect which requires non-parallel temperature and electron density gradients. While exploring the physics at a comparable epoch, we consider here a  different and more efficient process, as will appear below.  In the following section, we describe in details the mechanism of our model. Then, in Sect. 3, we present the results obtained for the cosmological case, emphasising on the magnetic field power spectrum and its order of magnitude at protogalactic and galactic scales. We also apply our mechanism to the case of  interstellar magnetic fields. Finally, we discuss the model and its limitations in Sect. 4. ",
        "conclusions": "We presented a new model for the generation of large scale magnetic fields. The mechanism proposed relies upon the simple physics of local charge separation which occurs naturally in the cosmological context of reionisation, at a redshift $z \\sim 7$. As we have seen, the first luminous objects, such as quasars, emit a radiation which ionises the gas filling the Universe. The anisotropy of the ionising flux leads to the generation of electrical fields. The distribution of the ionising sources together with the gas density inhomogeneities imply that the ionising flux is inhomogeneous. The fluctuations in the flux generate electric currents which in turn induce a magnetic field on the scale of the fluctuations.  With only a couple of simple initial assumptions, we could derive interesting properties for the generated magnetic fields. In particular, the fields produced in our model have amplitudes several orders of magnitude higher than in usual plasma physics magnetogenesis models. These fields are therefore suitable to serve as seeds on large scales for adiabatic collapse amplification up to $10^{-10}$ Gauss in galaxies. Such large values represent an improvement for further amplification by dynamo mechanisms since less time is  required in order to account for magnetic fields observed in high redshift galaxies.  Furthermore, relating the magnetic field power spectrum to the matter density power spectrum, we have shown that magnetic fields are generated in our model preferably on large scales, whereas small scale fields are strongly suppressed. This result is true at least in the non-saturated regime and we are currently exploring the saturated case. Such a property is highly interesting again with respect to dynamo amplification theories: the small scale fields are weak and therefore their amplification by dynamo represent a less crucial threat.  We showed also that, in the context of the ISM, our mechanism is negligible as a source of magnetic fields when compared to usual effects such as dynamo and turbulence.  As we mentioned earlier, Subramanian \\textit{et. al.}~\\cite{subra} and Gnedin \\textit{et. al.}~\\cite{gnedin} explored the generation of magnetic field seeds in the context of reionisation. In their model, the magnetic field source term is the Battery effect relying on charge separation. The latter is obtained when ionisation fronts either break out from protogalaxies or propagate past overdense regions of the Universe. Across an ionising front, the temperature varies sharply by several orders of magnitude. The driving physical process for charge separation is the thermal pressure of the gas, acting in the same way on electrons and ions. Since their mass is far smaller than the ion mass, electrons are more accelerated by the pressure force than ions and this provides charge separation. However, the mechanism we presented in this paper relies on a quite different process: The source of charge separation is not differential thermal pressure but differential radiation pressure. For a particle $i=e,p$, the acceleration due to thermal pressure is $a^\\text{t}_i \\propto m_i^{-1}$ whereas radiation pressure provides an acceleration $a^\\text{r}_i \\propto m_i^{-3}$. Radiation pressure is therefore more efficient by a factor of \\beq \\frac{a^\\text{r}_\\text{e}/a^\\text{r}_\\text{p}}{a^\\text{t}_\\text{e}/a^\\text{t}_\\text{p}} = \\left(\\frac{m_\\text{p}}{m_\\text{e}}\\right)^2 \\approx 3. ~10^6 \\eeq than thermal pressure for accelerating electrons and creating electrical fields. This explains why our model leads to orders of magnitude higher predictions than in \\cite{subra}. Thus, the mechanism we described in this paper is likely to be dominant in cosmological situations where strong anisotropic and inhomogeneous radiation pressure exists.  Further developments of the model presented in this paper,  relaxing our geometrical assumptions and addressing the magnetic field power spectrum  in the saturated regime will appear in a forthcoming paper.\\\\"
    },
    "0212/astro-ph0212422_arXiv.txt": {
        "abstract": "We have investigated the X-ray spectral properties of a collection of low-mass X-ray binaries (LMXBs) within a sample of 15 nearby early-type galaxies using proprietary and archival data from the {\\it Chandra X-ray Observatory}. We find that the spectrum of the sum of the sources in a given galaxy is remarkably similar from galaxy to galaxy when only sources with X-ray luminosities less than $10^{39}$ ergs s$^{-1}$ (0.3--10 keV) are considered. Fitting these lower luminosity sources in all galaxies simultaneously with a power law model led to a best-fit power law exponent of $\\Gamma = 1.56 \\pm 0.02$ (90\\% confidence), and using a thermal bremsstrahlung model yielded $kT_{brem} = 7.3 \\pm 0.3$ keV. This is the tightest constraint to date on the spectral properties of LMXBs in external galaxies. The spectral properties of the LMXBs do not vary with galactic radius out to three effective radii. There is also no apparent difference in the spectral properties of LMXBs that reside within globular clusters and those that do not. We demonstrate how the uniformity of the spectral properties of LMXBs can lead to more accurate determinations of the temperature and metallicity of the hot gas in galaxies that have comparable amounts of X-ray emission from hot gas and LMXBs.  Although few in number in any given galaxy, sources with luminosities of $1-2 \\times 10^{39}$ ergs s$^{-1}$ are present in 10 of the galaxies. The spectra of these luminous sources are softer than the spectra of the rest of the sources, and are consistent with the spectra of Galactic black hole X-ray binary candidates when they are in their very high state. The spatial distribution of these sources is much flatter than the optical light distribution, suggesting that a significant portion of them must reside within globular clusters. The simplest explanation of these sources is that they are $\\sim 10-15$ M$_{\\odot}$ black holes accreting near their Eddington limit. The spectra of these sources are very different than those of ultraluminous X-ray sources (ULXs) that have been found within spiral galaxies, suggesting that the two populations of X-ray luminous objects have different formation mechanisms. The number of sources with apparent luminosities above $2 \\times10^{39}$ ergs s$^{-1}$ when determined using the distance of the galaxy is equal to the number of expected background AGN and thus appear to not be associated with the galaxy, indicating that very luminous sources are absent or very rare in early-type galaxies. The lack of ULXs within elliptical galaxies strengthens the argument that ULXs are associated with recent star formation. ",
        "introduction": "\\label{sec:intro}  The {\\it Chandra X-ray Observatory} has made it possible to resolve dozens if not hundreds of individual X-ray point sources in nearby galaxies owing to its sub-arcsecond spatial resolution (Sarazin, Irwin, \\& Bregman 2000; Angelini, Loewenstein, \\& Mushotzky 2001; Kraft et al.\\ 2001; Bauer et al.\\ 2001; Soria \\& Wu 2002). While spiral galaxies contain a variety of types of X-ray point sources (high- and low-mass X-ray binaries, supernovae remnants), elliptical and S0 galaxies, as well as the bulges of spiral galaxies contain almost exclusively low-mass X-ray binaries (LMXBs). LMXBs are believed to be composed of a compact accreting primary (a neutron star or a black hole) and a low-mass main sequence or red giant secondary that is losing material to the primary as a result of Roche lobe overflow. Although LMXBs have been studied extensively in our Galaxy, resolving them and determining their spatial distribution and spectral characteristics in external galaxies has only recently become feasible.  \\begin{table*}[t] \\scriptsize \\caption{Sample of Galaxies} \\label{tab:sample} \\begin{center} \\begin{tabular}{lcccccccccc} \\multicolumn{10}{c}{} \\cr \\tableline \\tableline Galaxy & Galaxy &Type & Obs. & Distance & Semimajor & Semiminor & $N_H$ & Exposure & Luminosity \\cr Number &&&ID & (Mpc) & Axis ($^{\\prime\\prime}$) & Axis ($^{\\prime\\prime}$) & ($10^{20}$ cm$^{-2}$) & (seconds) & Limit (ergs s$^{-1}$)\\cr \\tableline \\phn\\phn1 & NGC~1291 & Sa & 795 & 8.9 & \\ldots  & \\ldots &2.12 & 37,406 & $2.0 \\times 10^{37}$  \\\\ \\phn\\phn2 & NGC~1316 & S0 & 2022 & 21.5 & 132.2\\tablenotemark{a} & 90.5\\tablenotemark{a} & 1.88 & 24,478 & $1.7 \\times 10^{38}$  \\\\ \\phn\\phn3 & NGC~1399 & E1 & 319 & 20.0 & 44.6\\tablenotemark{b} & 40.5\\tablenotemark{b} & 1.34 & 54,540 & $5.7 \\times 10^{37}$   \\\\ \\phn\\phn4 & NGC~1407 & E0 & 791 & 28.8 & 73.9\\tablenotemark{b} & 68.7\\tablenotemark{b} & 5.42 & 33,763 & $2.8 \\times 10^{38}$   \\\\ \\phn\\phn5 & NGC~1549 & E0 & 2077 & 19.7 & 51.0\\tablenotemark{b} & 44.7\\tablenotemark{b} & 1.46 & 21,892 & $1.8 \\times 10^{38}$    \\\\ \\phn\\phn6 & NGC~1553 & S0 & 783 & 18.5 & 78\\tablenotemark{c} & 51\\tablenotemark{c} & 1.50 & 16,361 & $1.6 \\times 10^{38}$    \\\\ \\phn\\phn7 & NGC~3115 & S0 & 2040 & 9.7 & 93\\tablenotemark{d} & 35\\tablenotemark{d} & 4.32 & 36,979 & $2.6 \\times 10^{37}$    \\\\ \\phn\\phn8 & NGC~3585 & E/S0 & 2078 & 20.0 & 56\\tablenotemark{e} & 28\\tablenotemark{e} & 5.57 & 33,706 & $1.7 \\times 10^{38}$   \\\\ \\phn\\phn9 & NGC~4374 & E1 & 803 & 18.4 &58.2\\tablenotemark{b} & 53.3\\tablenotemark{b} &  2.60 & 28,049 & $9.8 \\times 10^{37}$   \\\\ \\phn\\phn10 & NGC~4472 & E2 & 321 & 16.3 & 114.0\\tablenotemark{f} & 95.6\\tablenotemark{f} & 1.66 & 26,326 & $9.5 \\times 10^{37}$    \\\\ \\phn\\phn11 & NGC~4494 & E1 & 2079 & 17.1 & 49.3\\tablenotemark{b} & 42.2\\tablenotemark{b} & 1.52 & 18,254 & $1.3 \\times 10^{38}$    \\\\ \\phn\\phn12 & NGC~4636 & E/S0 & 323 & 14.7 & 117.0\\tablenotemark{f} & 85.6\\tablenotemark{f} & 1.81 & 42,686 & $5.2 \\times 10^{37}$   \\\\ \\phn\\phn13 & NGC~4649 & E2 & 785 & 16.8 & 82.0\\tablenotemark{f} & 66.2\\tablenotemark{f} & 2.20 & 18,401 & $1.4 \\times 10^{38}$    \\\\ \\phn\\phn14 & NGC~4697 & E6 & 784 & 11.8 & 97.3\\tablenotemark{b} & 60.2\\tablenotemark{b} & 2.12 & 39,063 & $3.1 \\times 10^{37}$    \\\\ \\phn\\phn15 & M31 & Sb & 309 & 0.76 & \\ldots & \\ldots & 6.66 & 5,089 & $1.2 \\times 10^{36}$    \\\\  \\tableline \\end{tabular} \\end{center} \\tablenotetext{a}{Caon, Capaccioli, \\& D'Onofrio (1994)} \\tablenotetext{b}{Goudfrooij et al.\\ (1994)} \\tablenotetext{c}{Kormendy (1984) and Jorgensen, Franx, \\& Kjaergaard (1995)} \\tablenotetext{d}{Capaccioli, Held, \\& Nieto (1987)} \\tablenotetext{e}{Ryden, Forbes, \\& Terlevich (2001)} \\tablenotetext{f}{Peletier et al.\\ (1990)}  \\end{table*}  Initial work on the LMXB population of early-type galaxies has led to several interesting results. Sarazin, Irwin, \\& Bregman (2000, 2001) detected 90 X-ray sources in a {\\it Chandra} observation of the elliptical galaxy NGC~4697, and found a break in the luminosity function of the sources at a luminosity of $\\sim3 \\times 10^{38}$ ergs s$^{-1}$, intriguingly close to the Eddington luminosity of an accreting 1.4 M$_{\\odot}$ neutron star. One interpretation of this break is that it represents a division between black hole X-ray binaries and neutron star X-ray binaries, with only black hole binaries more luminous than the break, and a mixture of neutron star and black hole binaries less luminous than the break. Such breaks have also been found in other early-type galaxies (Blanton, Sarazin, \\& Irwin 2001; Finoguenov \\& Jones 2002; Randall, Sarazin, \\& Irwin 2002; Kundu, Maccarone, \\& Zepf 2002). Recent studies have also found that a significant fraction (40\\%--70\\%) of the X-rays binaries reside in globular clusters of the host galaxy (Angelini et al.\\ 2001; Randall et al.\\ 2002; Kundu et al.\\ 2002), in marked contrast to the $\\sim10\\%$ of Galactic and M31 LMXBs that reside within globular clusters. In addition, when the spectra of all the resolved LMXBs of a given galaxy were added together and fit with a power law spectral model, an index ranging from 1.5--2.0 has been obtained (Sarazin et al.\\ 2001; Randall et al.\\ 2001; Kim \\& Fabbiano 2002; Irwin, Sarazin, \\& Bregman 2002).  Previous X-ray satellites lacked the needed spatial resolution and bandpass coverage to separate cleanly the LMXB component from the hot gas component in early-type galaxies. While recent progress has been made from the study of the luminosity functions and host environment of LMXBs in early-type galaxies, a detailed analysis of the spectral properties of LMXBs with a large sample of galaxies has not yet been attempted with {\\it Chandra}. With {\\it Chandra} we are now in a position to determine the spectral properties of individual bright sources in galaxies or add up the spectra of the fainter sources to determine the bulk spectral properties of the LMXBs.  A more thorough understanding of the X-ray spectral properties of LMXBs in early-type galaxies is critical for the study of the hot, X-ray--emitting gas within these systems. Although the LMXB contribution to the X-ray emission from gas-rich galaxies is negligible, this is not the case for galaxies with moderate to low amounts of X-ray--emitting gas. In these galaxies, the LMXB contribution can be the dominant X-ray emission mechanism. Quantifying the spectral properties of LMXBs in nearby galaxies where the LMXBs are resolvable will allow for a more accurate separation of the gaseous and stellar X-ray components in galaxies too distant for the individual LMXBs to be detected.  In this paper we use both proprietary and archival {\\it Chandra} data for 15 early-type systems (consisting of eight elliptical galaxies, two transitional E/S0 galaxies, three S0 galaxies, and two spiral bulges) to determine the spectral characteristics of LMXBs over a range of X-ray luminosity classes, as well as a function of galactic radius. Unless otherwise stated, all uncertainties are 90\\% confidence levels, and all X-ray luminosities are in the 0.3--10 keV energy band. ",
        "conclusions": "\\label{sec:conclusions}  We have used {\\it Chandra} data for 15 early-type systems to constrain the spectral and spatial properties of LMXBs in these galaxies. We have found that once the most luminous ($> 10^{39}$ ergs s$^{-1}$) sources are removed from the combined spectra of the sources, the spectra of the sum of the sources are very similar values among the galaxies. When all the galaxies are fit simultaneously with a power law spectral model, the best-fit power law exponent is $\\Gamma = 1.56 \\pm 0.02$. Even faint sources as dim as $10^{36}$ ergs s$^{-1}$ in the bulge of M31 have similar spectral properties as the more luminous sources. There was no apparent difference in the spectral properties of LMXBs as a function of galactic radius. Nor was there a significant difference in the spectral properties of sources based on their presence within or outside a globular cluster.  A significant number of sources with luminosities of $1-2 \\times 10^{39}$ ergs s$^{-1}$ were found within the galaxies, and they exhibited significantly softer spectral properties than the fainter sources. The disk blackbody + power law model used to model their spectra is very reminiscent of Galactic black hole X-ray binaries when they are in their very high state. Their spectra were also quite different from ULXs found within spiral galaxies. The simplest explanation of these sources is that they are $\\sim7-15$ M$_{\\odot}$ accreting near their Eddington limit. The spatial distribution of these sources is significantly more extended than the optical light.  With rare exception, sources more luminous than $2 \\times 10^{39}$ ergs s$^{-1}$ are absent from early-type galaxies. The number and spatial distribution of the sources with fluxes corresponding to $2 \\times 10^{39}$ ergs s$^{-1}$ or greater if they are at the distance of the galaxy is consistent with them being unrelated background/foreground sources. The only exceptions to this seems to be two $\\sim 4\\times 10^{39}$ ergs s$^{-1}$ sources found within globular clusters of NGC~1399. Their spectra are also quite different than that of a typical spiral galaxy ULX. Their presence within a globular cluster suggests that globular clusters might harbor intermediate mass black holes that are accreting at a few percent of their Eddington limit.  Finally, we have discussed how these constraints on the spectral properties of LMXBs, especially the fainter ones, can lead to better constraints to the luminosity, temperature, and metallicity of the hot gas within early-type galaxies that contain little gas."
    },
    "0212/astro-ph0212570_arXiv.txt": {
        "abstract": "{  Low cooling plasmas associated with large kinetic energies are likely to be the origin of the kpc-extended and well collimated extra-galactic jets.  It is proposed that jets are launched  from a layer, governed by a highly diffusive, super-Keplerian rotating  and thermally dominated by virial-hot and magnetized ion-plasma. The launching layer is located between the accretion disk and the corona surrounding the nucleus. The matter in the layer is causally connected to both the disk and to the central engine. Moreover we find that coronae, in the absence of heating from below,  are  dynamically unstable to thermal ion-conduction, and that accretion disks  become intrinsically advection-dominated.  We confirm the capability of this multi-layer model to form jets by carrying out 3D axisymmetric quasi-stationary MHD calculations with high spatial resolution, and taking into account turbulent and magnetic diffusion. The new multi-layer topology accommodates several previously proposed elements for jet-initiation, in particular the ion-torus,  the magneto-centrifugal and the truncated disk - advective tori models. ",
        "introduction": "\\label{intro} Jets have been observed in many systems including active galaxies, X-ray binaries, black holes X-ray transients, supersoft X-ray sources and young stellar objects (K\\\"onigl 1997, Livio 1999, Mirabel 2001). Each of these systems is considered to contain an accretion disk, while jet-speeds have been verified to be of the order of the escape velocity at the vicinity of the central object (Mirabel 1999, Livio 1999 and the references therein). Recent observations of the M87 galaxy reveal a significant jet-collimation already at 100 gravitational radii from the central engine, and that jet-launching should occur close to the last stable orbit (Biretta et al. 2002).  Several scenarios have been suggested to uncover the mechanisms underlying jet-initiations and their connection to accretion disks (Pudritz \\& Norman 1986). In most of these models magnetic fields (-MFs) are considered to play the major role in powering and collimating jets (e.g., the magneto-centrifugal acceleration model of Blandford \\& Payne 1982, the ion-torus model of Rees et al. 1982, X-point model of Shu et al. 1995, ADAF and ADIOS models of Narayan \\& Yi 1995 and Blandford \\& Begelman 1999).  Previous radiative hydrodynamical studies without magnetic fields have confirmed the formation of a transition layer (-TL) between the disk and the corona, governed by thermally-induced outflows (Hujeirat \\& Camenzind 2000). The aim of this paper is show that incorporating large scale magnetic fields (-MFs) manifests such formation and dramatically strengthen the dynamic of the in- and out-flows. Moreover, the TL is shown to be an optimal runaway region where highly energetic ion-jets start off. The back reaction of jet-flows on the structure of the disk and on the corona surrounding the nucleus is investigated also.  The study is based on self-consistent 3D axi-symmetric quasi-stationary MHD calculations, taking into account magnetic and hydro-turbulent diffusion, and adopting the two-temperature description (Shapiro et al. 1976). This adaptation is fundamental as 1) the dynamical time scale around the last stable orbit may become shorter than the Coulomb-coupling time. Therefore turbulent dissipation, adiabatic or shock compression preferentially heat up the ions rather than electrons ($T_\\mathrm{i} \\propto \\rho^{2/3}_\\mathrm{i}$, while $T_\\mathrm{e} \\propto \\rho^{1/3}_\\mathrm{i}$). 2) Taking into account that ions radiate inefficiently, having virial-heated ions in the vicinity of the last stable orbit is essential for the total energy-budget of large scale jets. \\begin{figure}[htb] \\centering {\\hspace*{-0.1cm} \\includegraphics*[width=8.0cm, bb=0 0 392 194,clip] {Eg221_f1.eps} } {\\vspace*{-0.1cm}} \\caption{ The model consists of  $10^8M_{\\odot}$ Schwarzschild BH at the center (its gravity is described in terms of the quasi--Newtonian potential of Paczynski\\,\\&\\,Wiita 1980), and an  SS-disk (blue color, extending from r=1 to r=20 in units of the radius of the last stable orbit i.e., in $3\\times R_\\mathrm{Schwarzschild}$, thickness $H_\\mathrm{d} = 0.1 r,$ an accretion rate of $\\Mdot = 0.01\\times \\Mdot_\\mathrm{Edd},$ and a central disk temperature of $T=10^{-3} T_\\mathrm{virial}$ at the outer radius). The ion-temperature $\\rm T_\\mathrm{i}$  is set to be equal to the electron temperature $\\rm T_\\mathrm{e}$ initially.) The low-density hot corona ($T=T_\\mathrm{virial}$, and density $\\rho(t=0,r,\\theta)=10^{-4}\\rho(t=0,r,\\theta=0$) is set to envelope the disk. A large scale magnetic field is set to thread the disk and the overlying corona (blue lines, $\\beta = P_\\mathrm{mag}/P_\\mathrm{gas} = 1/4$ at the outer radius, where $P_\\mathrm{gas}$ is the ion-pressure,  $P_\\mathrm{mag} = B\\cdot B/8 \\pi$ is the magnetic pressure, and $\\rm B$ is the magnetic field whose components are $(B_\\mathrm{P}, B_\\mathrm{T})= (B1,B2,B_\\mathrm{T})$.) The low-density hot corona ($T=T_\\mathrm{virial}$, and density $\\rho(t=0,r,\\theta)=10^{-4}\\rho(t=0,r,\\theta=0$) is set to envelope the disk. A large scale magnetic field is set to thread the disk and the overlying corona (blue lines, $\\beta = P_\\mathrm{mag}/P_\\mathrm{gas} = 1/4$ at the outer radius, where $P_\\mathrm{gas}$ is the ion-pressure,  $P_\\mathrm{mag} = B\\cdot B/8 \\pi$ is the magnetic pressure, and $\\rm B$ is the magnetic field whose components are $(B_\\mathrm{P}, B_\\mathrm{T})= (B1,B2,B_\\mathrm{T})$.) The numerical procedure is based on using the implicit solver IRMHD3  to search steady-state solution for the 3D axi-symmetric two-temperature diffusive MHD equations in spherical geometry (for further clarifications about the equations and the numerical method see Hujeirat \\& Rannacher 2001, Hujeirat \\& Camenzind 2000). The ion-pressure is used to describe the turbulent viscosity: $\\nu_\\mathrm{turb} = \\alpha P_\\mathrm{gas}/\\Omega$, where $\\alpha$ is the usual viscosity coefficient, and $\\Omega$ is the angular frequency. The magnetic diffusivity is taken to be equal to $\\nu_\\mathrm{turb}.$ $250\\times80$ strongly stretched finite volume cells in the radial and vertical direction have been used. Normal symmetry and anti-symmetry boundary conditions have been imposed along the equator and the rotation axis. Extrapolation has been adopted to fix down-stream values at the inner boundary. Non-dimensional formulation is adopted, using the reference scaling variables: $\\rm{\\tilde{\\rho}= 2.5 \\times 10^{-12} {\\rm g\\, cm^{-3}}, \\tilde{T}= 5 \\times 10^7 K},$ $\\tilde{U}=\\tilde{V_\\mathrm{S}} = \\gamma {\\cal R}_\\mathrm{g} \\tilde{T}/ \\mu_\\mathrm{i},$ $(\\mu_\\mathrm{i}=1.23)$. $\\tilde{B}=\\tilde{V_\\mathrm{S}} \\sqrt{4 \\pi \\tilde{\\rho}}.$ The location of the transition layer (-TL), where the ion-dominated plasma is expected to rotate super-Keplerian and being accelerated into jets, is shown for clarity. } \\end{figure} ",
        "conclusions": "We have presented a multi-layer model for initiating ion-jets in AGNs. The model agrees with a previous numerical study which confirmed the formation of thermally-induced outflows in the transition layer (Hujeirat \\& Camenzind 2000). Here we have shown that incorporating the effects of MFs manifests their formation and dramatically strengthen their dynamics.  Three ingredients for initiating winds have been detected: 1) a highly diffusive plasma dominated by virial-hot ions, 2) large scale magnetic fields that efficiently transport angular momentum from the disk into the TL, where the plasma rotates super-Keplerian, and 3) an underlying advection-dominated accretion disk.  Taking into account that the corona is dynamically unstable, adopting a large scale magnetic topology, and allowing ion-electron thermal decoupling appear to force accretion flows to undergo a global energy re-distribution: confined inflows (negative Bernoulli number) in the equatorial region and in the corona, and thermally and magneto-centrifugally-driven outflows in the TL characterized through a positive Bernoulli number. This feature may survive under different conditions: strong MFs suppress turbulence, weakening thereby the effect of the turbulent-viscosity and dominate the transport of angular momentum. On the other hand weak MFs in rotating stratified flows would be amplified via dynamo-action and reach equipartition, beyond which turbulence is again suppressed. This interplay between MFs and turbulent-viscosity, Balbus-Hawley and Parker instabilities  may settle into an equilibrium state, in which inflows are simultaneously associated with low-cooling out-flowing ion-plasma.  We note that in the absence of thermal conduction and adopting the one-temperature description,  low-viscosity radiatively inefficient HD and MHD accretion flows become inevitably convection-dominated. Therefore, in the early phases of jet-initiation, CDAFs may play an important role in powering the jets in AGNs and microquasars (Abramowicz et al. 2002).  The multi-layer model presented here accommodates some elements of BP82. In particular, we agree with BP82 about the necessity for a super-Keplerian rotation of the plasma overlying the accretion disk. However, the plasma here is dominated by highly-diffusive and virial-hot ions; it does not require a special $\\rm{B_\\mathrm{P}}-$alignment with respect to the disk-normal to enable jet-launching, as ideal-MHD treatment requires. While our results agrees with the ion-torus model with respect to the necessity of 2T-plasma to  maintain the ions hot for a significant time of their propagation-life in the ISM, no signatures for the formation of ion-supported tori have been detected (Rees et al. 1982).\\\\  Our results differ from ADAF and ADIOs in several issues, and in particular  with respect to 1) the existence of a layer adjusting to the disk, where the plasma is found to rotate super-Keplerian, 2) the configurations of the in- and the out-flows, 3) stability of the corona in the vicinity of the BH, 4) the transition from SS-disk to advection-dominated disks and 5) with respect to the essence of Bernoulli number in dissipative flows, i.e., a positive Bernoulli number is necessary but not sufficient for outflows (Abramowicz et al. 2000).\\\\  Finally, we note that since the flow in the TL is highly dissipative (strengthen thermal and rotational coupling with the central nucleus), the innermost region of the disk rotates synchronously with Kerr black holes (due to the frame dragging effect), and since $\\tau_\\mathrm{rem}$ decreases with radius and depends inversely on $\\rm{B_\\mathrm{T}}$ and $\\rm{B_\\mathrm{P}}$, we think that the plasma attached to the poloidal magnetic field would be forced to deposit its angular momentum to the plasma in the TL, thereby considerably enhancing the centrifugal power and ejecting the ion-plasma into space with relativistic speeds. \\small"
    },
    "0212/astro-ph0212020_arXiv.txt": {
        "abstract": "A technique of timescale analysis performed directly in the time domain has been developed recently. We have applied the technique to study rapid variabilities of hard X-rays from neutron star and black hole binaries, $\\gamma$-ray bursts and terrestrial $\\gamma$-ray flashes. The results indicate that the time domain method of spectral analysis is a powerful tool in revealing the underlying physics in high-energy processes in objects. ",
        "introduction": "The widely used Fourier analysis method in temporal analysis is to derive frequency spectra from a time series. After the Fourier transform of a light curve $x(t_k)$ \\begin{equation} X(f_j)=\\sum_k x(t_k) e^{-i2\\pi f_jk\\Delta t}, \\hspace{5mm} f_j=j/T \\end{equation} we can get the power density spectrum \\begin{equation} P_j=|X(f_j)|^2 \\end{equation} From two light curves, $x_1(t_k)$ and $x_2(t_k)$, observed simultaneously in two energy bands at times $t_k$, and their Fourier transforms $X_1(f_j)$ and $X_2(f_j)$, we can construct the cross spectrum \\begin{equation}C(f_j)=X_1^*(f_j)X_2(f_j) \\end{equation} and then the time lag spectrum \\begin{equation} \\Lambda (f_j)=\\arg [C(f_j)]/2\\pi f_j \\end{equation} and the coherence coefficient spectrum \\begin{equation} r(f)=\\frac{|<C(f)>|}{\\sqrt{<|X_1(f)|^2><|X_2(f)|^2>}} \\end{equation}  People usually take a Fourier period $1/f$ as a timescale and use the Fourier power spectrum to describe the distribution of variation amplitude at different timescales, the time lag spectrum to describe the distribution of the emission time difference between two energy bands at different timescales and the coherence coefficient spectrum to describe the distribution of degree of linear correlation between high and low energy processes at different timescales. But it is not correct to equate the Fourier period with the timescale; a Fourier component with a certain frequency $f$ of a light curve is not equal to the real process with the timescale $1/f$. For some important high-energy emission processes the Fourier spectra distort the timescale distribution of real physical processes seriously. It is needed to derive timescale spectra from observed light curves directly in the time domain without through Fourier transforms. ",
        "conclusions": ""
    },
    "0212/astro-ph0212216_arXiv.txt": {
        "abstract": "Weak gravitational lensing of background galaxies by intervening matter directly probes the mass distribution in the universe. This distribution, and its evolution at late times, is sensitive to both the dark energy, a negative pressure energy density component, and neutrino mass. We examine the potential of lensing experiments to measure features of both simultaneously. Focusing on the radial information contained in a future deep $4000$ square degree survey, we find that the expected ($1$-$\\sigma$) error on a neutrino mass is $0.1$ eV, if the dark energy parameters are allowed to vary.  The constraints on dark energy parameters are similarly restrictive, with errors on $w$ of $0.09$. Much of the restrictive power on the dark energy comes not from the evolution of the gravitational potential but rather from how distances vary as a function of redshift in different cosmologies. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/astro-ph0212166_arXiv.txt": {
        "abstract": "Gas processes affecting star formation are reviewed with an emphasis on gravitational and magnetic instabilities as a source of turbulence. Gravitational instabilities are pervasive in a multi-phase medium, even for sub-threshold column densities, suggesting that only an ISM with a pure-warm phase can stop star formation.  The instabilities generate turbulence, and this turbulence influences the structure and timing of star formation through its effect on the gas distribution and density.  The final trigger for star formation is usually direct compression by another star or cluster.  The star formation rate is apparently independent of the detailed mechanisms for star formation, and determined primarily by the total mass of gas in a dense form. If the density distribution function is a log-normal, as suggested by turbulence simulations, then this dense gas mass can be calculated and the star formation rate determined from first principles. The results suggest that only $10^{-4}$ of the ISM mass actively participates in the star formation process and that this fraction does so because its density is larger than $10^5$ cm$^{-3}$, at which point several key processes affecting dynamical equilibrium begin to break down. ",
        "introduction": "Gas processes in the interstellar medium (ISM) are varied and complex. This review is limited to those most closely involved with precursors to star formation. Other talks at this conference cover the high energy phase and the dispersal of gas after star formation.  Some ideas expressed here are considered in more detail in Elmegreen (2002).  At the beginning of star formation is cloud formation, but stars are also triggered in pre-existing clouds by processes unrelated to cloud formation (e.g., by supernovae), and many clouds are formed that do not produce new stars (e.g., diffuse clouds). Thus star formation is distinct from cloud formation.  Figure 1 shows a diagram of the flow of energy into ISM structure, starting with sources dominated by young stars, gaseous self-gravity, and magnetism (which derives its energy from galactic rotation via the dynamo).  The stellar sources tend to produce expanding regions and cosmic rays, turning their energy into radiation behind shock fronts and turbulent motions that also decay into radiation.  Gravity produces contracting regions by swing amplified instabilities and collapse along spiral arms. This contraction releases more gravitational energy as the density increases, and again much of this energy goes into turbulence and ultimately radiation. The shells produced by stellar pressures and the turbulence produced in these shells and by various instabilities makes the observed cloudy structure of the ISM. Other stellar pressures, along with continued self-gravity and magnetic forces, then modify these clouds and eventually produce individual and binary stars on very small scales. \\begin{figure} \\centerline{\\includegraphics[width=28pc]{elmegreen.fig1.eps}} \\caption{Schematic diagram showing paths from the main energy sources, which are self-gravity, magnetic fields, and stars, to the formation of cloudy structure, going through intermediate steps of explosions, instabilities, and turbulence.  The cloudy structure that is formed by these processes is modified further by stars, gravity and magnetic fields.} \\end{figure} ",
        "conclusions": "Instabilities involving gravity, magnetism, and pressure lead to spirals, accretion, clouds, and turbulence.  Stellar pressures produce bubbles, more turbulence, and triggered star formation in clouds that already formed. Self-gravity and turbulence combine to structure the ISM, giving self-correlated properties for the gas and young stars with respect to size, velocity dispersion, and crossing time or duration of star formation. Turbulence also gives power law mass functions for clouds and clusters.  The turbulence generated by gravitational instabilities can maintain the ISM in a state of quasi-equilibrium where $\\Sigma\\sim\\Sigma_{\\rm crit}$.  If small scale instabilities continue in the cool component of the gas even when the average rms speed is large enough to give global stability, then star formation cannot regulate the $\\Sigma/\\Sigma_{\\rm crit}>1$ threshold. In this case, there is no self-regulation of star formation involving $\\Sigma_{\\rm crit}$ on a galactic scale. This will be true even if young stellar pressures agitate the ISM locally. They can blow the gas out into the halo and stop star formation locally, but young stars probably cannot fine-tune or moderate their own formation rate so that it stays near the historical or Hubble-type average. Young stars commonly trigger other stars anyway, so the feedback they produce should de-stabilize, not stabilize, the star formation rate, unless the entire local ISM is removed.  The star formation rate depends on the mass fraction in dense gas. Turbulence may determine this mass fraction, independent of the sources for the turbulence. The global SFR is then independent of the detailed triggering mechanisms. Then again there would be no self-regulation of star formation, only a star formation {\\it saturated} to its maximum possible value, as determined by the open and tenuous geometry of the gas.  In this case, star formation can be halted only by a dominance of the warm phase of the ISM."
    },
    "0212/astro-ph0212485_arXiv.txt": {
        "abstract": "s{ I discuss anew how arguments about the internal dynamics of galactic disks set constraints on the otherwise ambiguous decomposition of the rotation curves of spiral galaxies into the contributions by the various constituents of the galaxies. Analyzing the two sample galaxies NGC\\,3198 and NGC\\,2985 I conclude from the multiplicities of the spiral arms and the values of the $Q$ disk stability parameters that the disks of both galaxies are `maximum disks'.} ",
        "introduction": "The rotation curves of spiral galaxies provide the most direct evidence for the presence of dark matter in galaxies. However, taken alone they do not discriminate luminous from dark matter because their decomposition into the contributions from the various constituents of the galaxies is highly ambiguous. Thus further constraints are needed. Considerations of the dynamical state of the resulting disk models can provide such constraints (Bosma 1999, Fuchs 1999). Of particular interest is the question if the much discussed `maximum--disk' models, i.e.~disks with their masses chosen at the maximum allowed by the data, are dynamical viable disk models. The degeneracy of the decomposition problem is illustrated in Fig.~1 for the example of NGC\\,3198. The rotation curve is modelled as \\begin{equation} v_{\\rm c}^2(R) = v_{\\rm c, disk}^2(R) + v_{\\rm c, halo}^2(R) + v_{\\rm c, is\\,gas}^2(R)\\,, \\end{equation} where $v_{\\rm c, disk}$, $v_{\\rm c, halo}$, and $v_{\\rm c, is\\,gas}$ denote the contributions due to the stellar disk, the dark halo, and the interstellar gas, respectively. The disk is modelled as an exponential disk and the dark halo is described by a quasi--isothermal sphere. In the left panel of Fig.~1 the maximum--disk model of Broeils (1992) is reproduced, while the right panel shows a submaximal disk model. The halo model parameters have been changed so that both fits to the observed rotation curve are of the same quality.  \\begin{center} \\hbox{ \\epsfclipon \\epsfxsize=5.4cm \\epsffile{york_f1.eps} \\hspace{0.2cm} \\epsfclipon \\epsfxsize=5.4cm \\epsffile{york_f2.eps} } \\end{center} \\vspace{- 0.5cm} {\\footnotesize Figure 1. Maximum and submaximal disk models of the rotation curve of NGC\\,3198.} ",
        "conclusions": "Arguments for and against maximum disks have been discussed at length in the literature and are reviewed in detail, for instance, by Bosma (1999) or Sellwood (1999). One of the major consequences of maximum disks are the implied large core radii of the dark halos. These challenge the contemporary theory of the formation of galaxies according to CDM cosmology.  Obviously the dynamical constraints on the decomposition of rotation curves must be tried out on a much larger data set than here in order to test the maximum disk hypothesis. This can be easily done with the density wave theory criterion by inspecting images of galaxies for which rotation curves are available. I have, for instance, analyzed a set of low surface brightness galaxies and found again indications for maximum disks (Fuchs 2003a, b). However, velocity dispersions have been measured in very few galaxies. Thus the $Q$ stability parameter criterion, which is an independent consistency check on the density wave theory criterion, can be applied only in deplorably few cases."
    },
    "0212/gr-qc0212121_arXiv.txt": {
        "abstract": "Gravitational physics of VLBI experiment conducted on September 8, 2002 and dedicated to measure the speed of gravity (a fundamental constant in the Einstein equations) is treated in the first post-Newtonian approximation. Explicit speed-of-gravity parameterization is introduced to the Einstein equations to single out the retardation effect caused by the finite speed of gravity in the relativistic time delay of light, passing through the variable gravitational field of the solar system. The speed-of-gravity 1.5 post-Newtonian correction to the Shapiro time delay is derived and compared with our previous result obtained by making use of the post-Minkowskian approximation. We confirm that the 1.5 post-Newtonian correction to the Shapiro delay depends on the speed of gravity $c_g$ that is a directly measurable parameter in the VLBI experiment. ",
        "introduction": "The gravitational VLBI experiment for measuring the relativistic effect of propagation of gravity field of orbiting Jupiter was conducted on September 8, 2002 by the National Radio Astronomical Observatory (USA) and the Max Plank Institute for Radio Astronomy (Germany). The idea of the experiment was proposed by Kopeikin \\cite{2} who had noted that an arbitrary-moving gravitating body deflects light not instantaneously but with retardation caused by the finite speed of gravity propagating from the body to the light ray. This is because positions of the gravitating bodies are connected to the light-ray particle by the gravity null cone which is defined by the equation for the retarded time in the gravitational {\\it Li\\'enard-Wiechert}  potentials being solutions of the Einstein equations.  The goal of the experiment was to confirm this prediction by making use of the close celestial alignment of Jupiter and the quasar J0842+1835. The goal of this paper is to make use of the parameterized post-Newtonian technique to reveal how the gravity propagation affects the light-ray trajectory. This approach allows us to confirm our previous theoretical results \\cite{2} by looking at the problem from different position.  The problem is formulated as follows. Light ray is emitted by a quasar (QSO J0842+1835) at the time $t_0$ and moves to the network of VLBI stations located on the Earth. As light moves, it passes through the variable gravitational field of the solar system, created by the orbital motion of Sun and other massive planets, and is received by the first and second VLBI stations at the times $t_1$ and $t_2$ respectively. Gravitational field of the solar system causes delay in the time of propagation of radio signals -- the effect discovered by \\cite{6} in static gravitational field approximation. We noted \\cite{2} that the present-day accuracy of phase-reference VLBI measurements is sufficient to detect a 1.5 post-Newtonian\\footnote{The first 1.0 post-Newtonian approximation is of order $c^{-2}$, and 1.5 post-Newtonian correction to the Newtonian gravity is of order $c^{-3}$. } correction to the Shapiro time delay incorporated implicitly into position of the light-ray deflecting body (see Fig.\\ref{fig1}) which must be taken at the retarded instant of time related to the time of observation by the equation of the gravity null-cone. We solved the Einstein-Maxwell system of differential equations and calculated this correction for Jupiter by making use of the post-Minkowskian approximation \\cite{2} with the retarded {\\it Li\\'enard-Wiechert} solutions of the Einstein equations. The retardation is caused by the finite speed of propagation of gravity $c_g=c$ which, as we have revealed in \\cite{2}, enters the 1.5 post-Newtonian correction to the Shapiro time delay. Hence, we concluded that the VLBI experiment on September 8, 2002 is sensitive to the effect of the {\\it propagation of gravity} and can be used to measure its speed.  Our primary paper \\cite{2} has given no particular details on the parameterization of the gravity propagation effect the amplitude of which was characterized by the parameter $\\delta$. For this reason the physical meaning of the gravity propagation effect was not understood adequately by some researchers \\cite{1,w-astro}. By scrutiny inspection we found that \\cite{1} made an erroneous use of insufficient 1.0 post-Newtonian (static) approximation to discuss the high-order 1.5 post-Newtonian (dynamic) correction to the Shapiro time delay while \\cite{w-astro} gave analysis of the problem of propagation of light ray in the time-dependent gravitational field of the solar system on the basis of the PPN formalism which was not sufficiently elaborated to tackle properly the 1.5 post-Newtonian approximation to general relativity (see appendix \\ref{a2}).  In this paper we extend our theoretical analysis of the VLBI experiment by proving that the parameter $\\delta=c/c_g-1$, where $c$ is the speed of light and $c_g$ is the speed of gravity. Parameter $\\delta$ is a measurable quantity and its determination will provided us with the numerical estimate of the magnitude of the propagation of gravity effect in terms of the speed of gravity $c_g$. The spectrum of plausible values of $c_g$ ranges from $c_g=c$ in general relativity to $c_g=\\infty$ as advocated by \\cite{tvf}. The VLBI experiment gives $c_g/c=1.06\\pm 0.21$ \\cite{apj}. The values $c_g<c$ however are unlikely and ruled out, for example, by observations of cosmic rays \\cite{mn}.  The paper is organized as follows. In section 2 we discuss the linearized Einstein equations and the speed-of-gravity parameterization that is used in section 3 for derivation of the post-Newtonian metric tensor. Section 4 deals with the post-Newtonian time delay formula for light propagation. The effect of retardation of gravity in the gravitational time delay is explained in section 5. The differential VLBI time delay is derived in section 6 where we prove that our previous post-Minkowskian formulation of the VLBI delay is identical with the post-Newtonian presentation. Section 7 contains discussion of the basic results of the present paper.  Roman indices run from 1 to 3 while Greek indices run from 0 to 3. Repeated indices assume the Einstein summation rule. Indices are raised and lowered with the Minkowski metric $\\eta_{\\alpha\\beta}={\\rm diag}(-1,+1,+1,+1)$. The Kroneker symbol $\\delta^\\alpha_\\beta={\\rm diag}(1,1,1,1)$. We denote coordinates $x^\\alpha=(ct,x^i)$ and partial derivatives $\\partial_\\alpha\\equiv(c^{-1}\\partial/\\partial t, \\partial/\\partial x^i)$. Partial derivative associated with $c_g$ is denoted as $\\eth_\\alpha\\equiv(c_g^{-1}\\partial/\\partial\\tau, \\partial/\\partial x^i)$, where $\\tau$ relates to time $t$ via Eq. (\\ref{1}).  We use the boldface italic letters to denote spatial vectors, e.g, ${\\bm a}\\equiv a^i=(a^1,a^2,a^3)$. The dot between two spatial vectors is the Euclidean scalar product, ${\\bm a}{\\bm\\cdot} {\\bm b}=a^1b^1+a^2b^2+a^3b^3$. The cross between two spatial vectors denotes the Euclidean vector product, ${\\bm a}\\times{\\bm b}\\equiv \\varepsilon_{ijk}a^jb^k$, where $\\varepsilon_{ijk}$ is the Levi-Civita symbol such that $\\varepsilon_{123}=+1$. ",
        "conclusions": "We have constructed a speed-of-gravity parameterization of the linearized equations of general relativity and calculated a speed-of-gravity correction to the Shapiro time delay given by Eq. (\\ref{6bb}). Its post-Newtonian expansion is identical with Eq. (12) from our previous work \\cite{2} where we used the post-Minkowskian approach and introduced the fitting parameter $\\delta$ on a phenomenological basis. The present paper proves that on the sequence of the parametric space-times manifolds, $\\delta=c/c_g-1$ and distinguishes two fundamental speeds in the Einstein and Maxwell equations. Hence, the relativistic VLBI experiment conducted on September 8, 2002 allows to determine the speed of gravity and to measure its numerical value.  Einstein theory of general relativity predicts that perturbation of the gravitational field produced by orbital motion of Jupiter propagates with the speed of gravity $c_g=c$ towards the position of the light-ray particle traveling from quasar to Earth. This prediction can be tested by VLBI technique without direct detection of gravitational waves. The VLBI experiment gives the magnitude of the retardation-of-gravity effect in terms of its speed $c_g/c=1.06\\pm 0.21$ \\cite{apj}. We conclude that the speed of gravity is the same as the speed of light within our experimental error.  \\subsection"
    },
    "0212/astro-ph0212350_arXiv.txt": {
        "abstract": "We present the results of a VLT and HST imaging survey aimed at the identification of $4.5<z<6$ galaxies. In the VLT data, a set of broad and intermediate band filters has been used to select 13 high--$z$ candidates in a $Z_{AB}\\leq 25$ {\\it mags} catalog, over an area of about 30 arcmin$^2$. Discrimination against lower redshift interlopers (mainly early--type galaxies at high redshift and cool Galactic stars) has been done combining morphological and spectral classification.  This sample has been combined with a deeper $I_{AB}\\leq 27.2$  {\\it mags} sample obtained from the Hubble Deep Field campaigns.  The VLT final sample consists of 13 high--$z$ candidates, 4 of which are identified with high confidence as $z>4.5$ galaxies. The resulting integral surface density of the $Z_{AB}<25$ candidates at $z>4.5$ is in the range $0.13-0.44$ arcmin$^{-2}$ and that in the highest redshift bin $5<z\\leq 6$ is between $0.07-0.13$ arcmin$^{-2}$.  In the two HDFs, we identify at $I_{AB}\\leq 27.2$ 25 galaxies in the range $4.5\\leq z < 5$ and 16 at $5\\leq z\\leq 6$, corresponding to surface densities of $3.1$ arcmin$^{-2}$ and $2$ arcmin$^{-2}$, respectively.  We show that the {\\it observed} $Z_{AB}<25$ UV luminosity density appear to drop by about one order of magnitude from $z\\simeq 3$ to $z\\simeq 6$.  However, if we apply a threshold to obtain an absolute--magnitude limited sample, the UV luminosity density results to be roughly constant up to $z\\simeq 6$.  We finally show that recent semi-analytic hierarchical models for galaxy formation, while predicting a nearly constant {\\it total} UV luminosity density up to $z\\simeq 6$, under-predict the observed UV luminosity density at $Z_{AB}\\leq 25$ and over-predicts the $I_{AB}\\leq 27.2$ one. This behaviour can be understood in term of a poor match to the slope of the UV luminosity function. ",
        "introduction": "Current findings on the production of the stellar baryon budget of the Universe (Madau et al. 1998), have shown a complex, gradual process, spread across a considerable fraction of the cosmological lifetime, rather than confined into a preferred epoch in the past.  The evolution of the rest-frame ultraviolet luminosity density can be used to trace this process: since it is produced mainly by O, B short-living massive stars, it is almost independent of the star formation history of the galaxies (Madau, Pozzetti \\& Dickinson 1998), but it is tied to the fraction of the current star formation rate not extincted by dust.  Measurements of the UV luminosity density from local environment up to $z\\simeq1$ show an increase with redshift (Lilly et al. 1996, Madau 1996,  Wilson et al 2002).  At the highest redshifts, where the results are often based on imaging surveys and photometric redshifts, the UV luminosity density  appears to be roughly constant between $z=1$ and $z\\sim4$ ( Connolly et al 1997, Pascarelle et al 1998, Fontana et al. 1999, Steidel et al 1999, Lanzetta et al. 2002)   Here we focus upon the redshift interval $z=4.5-6$, a range which marks a critical phase in the history of the Universe, when it was just emerging from the reionization epoch.  Although first glances of the early Universe have been made possible from Sloan Survey, that have provided the first statistically defined sample of quasars extending up to $z=6.28$, the galaxy population at these epochs is presently nearly unexplored.   Indeed it is at these very high redshifts that it is possible to set severe constraints on the physical mechanisms that drive the galaxy formation and evolution. The main goal behind the study of high redshift galaxies is the construction of a coherent picture of the physical processes that led to galaxy evolution.  Within the framework of gravitational instability driven by primordial fluctuation, simple but physically motivated prescriptions have been used over the last years to describe the main processes involved (e.g. White \\& Frenk 1991; Cole et al. 1994,2000; Somerville \\& Primack 1999; Menci et al. 2002). These models are not as ``tunable'' as commonly believed since they can not be modified freely to match the high redshift observables without worsening the local fits (e.g. the local luminosity function or local Tully-Fisher relation). In this respect, any discrepancy between the model predictions and the observed properties reveals the lack or incorrect treatment of some fundamental processes.  In this context, it is important to perform these comparisons directly at the highest redshifts, where the physical processes leading to present-day galaxies are caught in act.  Despite the obvious interest in the field, and the exciting results obtained at $z\\simeq 3$ (Steidel et al. 1996), the discovery of galaxies at $z>4.5$ has been so far serendipitous.  There are obvious difficulties to tackle: first, the objects become progressively fainter and rarer. Second, the multicolor ``drop--out'' criteria used to select high--$z$ galaxies must be shifted from the UV-visual into the near-IR bands to follow the rest--frame UV.  For the same reasons, the spectroscopic follow--up is progressively harder. Besides this, the number of possible interlopers increases, with early--type galaxies and late--type stars progressively entering in the selection criteria. The current statistics of high--$z$ galaxies reflects these problems.  Few serendipitous objects have been identified (e.g. Spinrad et al 1998), mostly because of a large EW emission line identified as Ly$\\alpha$, (e.g. Chen et al 1999; Hu et al. 1998, 1999, 2002). Some identification have been later disputed on the basis of deep imaging shortward of the (presumed) Lyman Limit (Chen et al 2000, Stern et al 2000).  In general, the search of $z\\geq5$ emission line galaxies has been shown to suffer from the strong contamination by OII emitters at $z\\simeq 1.4$ (Stern et al 2000).  Thus, the conclusions that can be drawn from the few objects observed so far are that {\\it a)} spectroscopy is by itself not conclusive at these redshifts to derive a firm estimate of the average cosmic star formation rate, firstly because it mostly requires the existence of a strong Ly$\\alpha$ emission line, and because of the contamination from OII emitters; and {\\it b)} deep imaging observations are required in any case not only to select the objects but also to validate the spectroscopic identifications.  Previous color selections at $z>4.5$ based on the photometric redshift technique (Fontana et al. 1999, 2002) took advantage of the very deep HST North and South samples, discovering 4 $z>5.5$ candidates in each field down to $I_{AB}<27.5$. Evidence has come up for a comoving UV luminosity density at $z\\simeq 5$ lower by a factor of 5 than at $z\\simeq3$.  However, this estimate is based on a sample of very faint sources selected in a small area and could be strongly affected by the presence of large-scale structure in the ``pencil'' beams.  For this reason, we have started a relatively deep survey of galaxies selected in the Z band to search for galaxies at $4.5 \\leq z \\leq 6.2$, covering two well known fields where broad band multicolor imaging is available, namely the extended ESO Imaging Survey (EIS) Hubble Deep Field South and the NTT Deep Field.  We have designed a strategy based on a combination of intermediate and broad band filters to identify high--$z$ galaxies against the increasingly large number of interlopers.  We emphasize that this is not a ``pre--selection'' survey, but it is designed to provide by itself a reliable identification for the bulk of the galaxy population at $z\\geq 5$.  To increase the accuracy of the photometric estimate of the redshift and to minimize the presence of interlopers (mainly intermediate redshift early--type galaxies and late type stars) in our high $z$ sample, we have included a set of intermediate band filters namely IB691, IB834 and IB915, available at the ESO/VLT FORS imager. These filters, although of intermediate width, are relatively efficient since they sample spectral regions devoid of strong sky emission lines. The adopted filter set is thus tailored to trace the peculiar spectral features of the extremely high redshift galaxies, i.e. the flat rest--frame UV continuum that even at $z\\sim 6$ can be sampled by the Z-IB915 color and the very abrupt change of the shape due to the hydrogen absorption spectral breaks by the intergalactic and interstellar medium sampled by the bluer bands.  In the following, we will adopt a $\\Lambda$CDM cosmology with $\\Omega_\\Lambda=0.7$, $\\Omega=1$ and $H_0=70$ km/s/Mpc. All magnitudes will be given in the AB system. ",
        "conclusions": "In this work we have presented the results of a pilot survey aimed at detecting $z>4.5$ galaxies with deep multicolor images. The key features of our approach are {\\it a)} the use of an extended set of broad (UBVRIZJK) and intermediate band ($\\Delta \\lambda \\simeq 400$~\\AA) filters centered at 6900 and 8340~\\AA, {\\it b)} the adoption of conservative thresholds on the S/N for object detection and {\\it c)} the application morphological and spectral criteria: all these aspects improve the estimates of the redshifts, and help exclude or minimize the number of lower $z$ interlopers. When applied to the present data set, these criteria make it possible to reject all the brightest interlopers, that would dominate the LD, and to select a ``minimal'' sample of 4 high--confidence candidates at $z>4.5$, and a sample of 9 additional candidates where the stellar contamination is uncertain.  This ambiguity depends on the relative depth of our data set: it is well within the possibilities of red--enhanced imagers at 8--m class telescope or of ACS to extend this kind of analysis 1-2 magnitudes fainter, so that the stellar contamination to the LD can be lowered by a large amount, as in the case of deeper but smaller WFPC2-HDF data that we present here.   Even with this ambiguity, the two selected datasets constrain the $z>4.5$ UV LD with sufficient accuracy to show that the {\\it observed} $Z_{AB}<25$ UV LD drops by about one order of magnitude from $z\\simeq 3$ to $z\\simeq 6$.  This drop is largely due to the progressively brighter cutoff in the rest frame luminosity function: if we correct for this incompleteness, the UV LD appears to be roughly constant from $z \\simeq 2.5$ up to $z\\simeq 6$.  The present results are apparently in contrast with recent findings by Lanzetta et al (2002), who make use of the same WFPC2--HDF data and apply photometric redshifts, that claim the global SFR to steadily increase up to $z=12$.  Unfortunately, the overall approach and techniques adopted are so different that it is difficult to make a clean comparison. First, we note that the use of wide aperture magnitude and the adoption of a much higher S/N threshold to the catalog strongly reduce the systematics in the WFPC2 data induced by surface brightness dimming, that otherwise require the complex and model--dependent correction applied by Lanzetta et al (2002). Besides, we draw our conclusions from a homogeneous sample, that includes only objects selected down to the same absolute magnitude limit, rather than correcting for the incomplete coverage of the luminosity function, as done by Lanzetta et al (2002).  Finally, the photometric redshift distribution obtained by Lanzetta et al (2002) is markedly different from our own, and peaks at $z\\simeq 0$, a factor that may lead to an overestimate of the faint end of the SFR distribution function that is used to correct the high--$z$ data. Despite all these differences, it is to be noted that the two results are still consistent when the more conservative estimate of Lanzetta et al (2002) is compared with our data in the appropriate redshift range $z\\simeq 3$ to $z\\simeq 6$: given the overall uncertainties both analyses suggest that UV LD is roughly constant in this redshift range.  The question that naturally arises is whether a constant UV LD up to $z\\simeq 6$ may be compatible with the hierarchical scenario of galaxy formation. To discuss this point, we present in both panels of fig. 3 the UV LD predicted by the CDM semianalytic model described in Menci et al. (2002). This model is based on the Cole et al (2000) recipes, with an additional improved treatment of aggregation of satellite galaxies in common DM haloes.  We first show that the {\\it total} SFR derived from this model (thin solid lines) is nearly constant from $z\\simeq 2$ to about $z\\simeq 5$, and then fades by only a factor about 5 at $z=6$.  To allow a fair comparison with our data, we have applied the effects of magnitude limit cut and dust extinction directly to the theoretical model. The effect is shown by the thick solid lines in Fig3a and Fig3b. The shaded area show the possible effects of dust extinction, ranging from no dust (upper lines) to SMC-like extinction curve (lower lines). Please note that the inclusion of dust correction effects in the theoretical model decreases the average UV LD respect to the unextincted amount.  It appears that while the CDM model broadly encompasses the observed values, it progressively under-predicts the LD observed in the bright $Z_{AB}<25$ sample, while over-predicts the LD observed in the faintest $I_{AB}<27.2$ sample. This behavior can be understood in term of a poor match to the slope of the UV luminosity function. At the brightest magnitudes, the CDM model under-predicts the number of bright sources, while it over-predicts the number of fainter sources dominating the deeper sample. This is analogous to what already found at $z\\simeq 3$ and confirms a general trend of this version of CDM models to under-predict the amount of star--formation rate in high redshift massive objects (Somerville and Primack 2001, Poli et al 2001 and Menci et al 2002, Cimatti et al 2002).  The origin of this discrepancy is likely tied to the lack or to the oversimplified treatment of fundamental physical process. Indeed, we remark that the basic recipes adopted in this model already concur to enhance the SFR in high redshift massive objects, and that these aspects are further boosted with respect to the models that we used in Fontana et al (1999). The star formation rate is computed as $\\dot M_*=M_g/\\tau_*$, where $M_g$ is the amount of cool gas and the time-scale $\\tau_*$ is proportional to the dynamical time $\\tau_{dyn}$ and to the galaxy circular velocity $V_c$ as $\\tau_*=\\epsilon_*^{-1}\\,\\tau_{dyn}\\,\\big(V_c/200\\,{\\rm km/s}\\big)^{\\alpha_*}$. Since both $M_g$ and $1/\\tau_{dyn}$ increase with redshift, and $\\alpha_*=-1.5$, these models naturally predict that the star-formation rate increases with redshift (for a given galaxy mass), and is more efficient in massive galaxies (at a given redshift). In addition, feedback effects are also a strong function of $V_c$, since they scale as $\\big(V_c\\big)^{-5.5}$, and again strongly favor the more massive objects. These basic recipes, combined with the effects of the biased process of galaxy formation induced by hierarchical merging, already conspire to boost at high $z$ the SFR in massive objects, with respect to the less massive one.  In the present context, it is not possible to flatten the high--$z$ luminosity function by simply changing the free parameters of the model, since one rapidly worsens the fit to the local observables (Cole et al 2000).  Other physical processes that are important in the high--$z$ Universe are not included in our model. On the one hand, mechanical and ionizational feedback on the intergalactic medium by early galaxy and QSO formation is not included in our rendition, a process that is expected to quench the SFR in low mass objects and hereby to flatten the low luminosity side of the LF. On the other hand, molecular cooling is efficient at high $z$ and not included here. Another possibility that has been proposed is that the starburst efficiency increases during major mergers.  These ``starburst'' models are known to increase the bright side of the LF at $z=3-4$ (Somerville and Primack 1999), but require new additional free parameters to be introduced in the models, and may become less efficient at $z>5$, when major merging events are very rare. The challenge of the next years is to include all these processes in a self consistent picture that reproduce the high--$z$ observables."
    },
    "0212/astro-ph0212399_arXiv.txt": {
        "abstract": "The current status of our efforts to trace cosmic structure with $10^6$ galaxies (2MASS), $10^3$ galaxy clusters (NORAS\\,II cluster survey), and precision measurements for $10^2$ galaxy clusters (\\gcss) is given. The latter is illustrated in more detail with results on the gas temperature and metal abundance structure for $10^0$ cluster (A1644) obtained with XMM-Newton. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/nucl-th0212065_arXiv.txt": {
        "abstract": "In this paper an equation of state of neutron star matter which includes strange baryons in the framework of Zimanyi and Moszkowski (ZM) model has been obtained. We concentrate on the effects of the isospin dependence of the equation of state constructing for the appropriate choices of parameters the hyperons star model. Numerous neutron star models show that the  appearance of hyperons is  connected with the increasing density in neutron star interiors.  Various studies have indicated that the inclusion of $\\delta$ meson mainly affects the symmetry energy and through this the chemical composition of a neutron star. As the effective nucleon mass contributes to hadron chemical potentials it alters the chemical composition of the star. In the result the obtained model of the star not only excludes large population of hadrons but also does not reduce significantly lepton contents in the star interior. ",
        "introduction": "Compact stars which are known from observations can be classified into two distinctive groups. The first one is exemplified by white dwarfs and the second by neutron stars. Neutron stars themselves are identify with pulsars and compact X-ray sources \\cite{weber}. At the core of a neutron star the density of matter ranges from a few times of the density of normal nuclear matter to about an order of magnitude higher. Thus various exotic forms of matter such as hyperons or quark-hadron mix phase are expected to emerge in the interior of a neutron star \\cite{weber}. The appearance of these additional degrees of freedom and their impact not only on neutron stars but on proto-neutron stars structure and evolution as well has been the subject of extensive studies \\cite{glen}. Properties of matter at such extreme densities are of particular importance in determining forms of equations of state relevant for neutron stars and in examining their global parameters \\cite{bema}. Theoretical description of hadronic systems should be performed with the use of quantum chromodynamics (QCD) as it is the fundamental theory of strong interactions. However, at the hadronic energy scale where the observed experimentally degrees of freedom are not quarks but hadrons  the direct description of nuclei in terms of QCD become inadequate. Other alternative approach has to be formulated one of which is quantum hadrodynamics (QHD) \\cite{walecka} giving quantitative description of the nuclear many body problem. QHD is a relativistic quantum field theory in which nuclear matter description in terms of baryons and mesons is provided. The original model (QHD-I) contains nucleons interacting through the exchange of simulating medium range attraction $\\sigma$ meson and $\\omega$ meson responsible for short range repulsion. Extension (QHD-II) of this theory \\cite{bog77,bodmer,gmuca} includes also the isovector meson $\\rho $. Nonlinear terms in the scalar and vector fields were added in order to get the correct value of the compressibility of nuclear matter and the proper density dependence in the vector self-energy. The variation of nucleon properties in nuclear medium is the key problem in nuclear physics. In order to incorporate quark degrees of freedom in the analysis of nuclear many-body system Guichon \\cite{qmc1} provides a quark-meson coupling model (QMC).  The extension of the QMC theory namely the quark mean field model (QMF), describing a nucleon with the use of the constituent quark  model, has been successfully applied to study the properties of both nuclear matter and finite nuclei. The model considered in this paper is an alternative version of the Walecka approach with enlarged meson sector. In the interior of neutron stars the density of matter could exceed normal nuclear matter density up to a few times, in such high density regime nucleon Fermi energies exceed the value of hyperon masses and thus the new hadronic degrees of freedom are expected to emerge. The higher the density the more various hadronic species are expected to populate. The onset of hyperon formation depends on the hyperon-nucleon and hyperon-hyperon interactions. Hyperons can be formed both in leptonic and baryonic processes. Several relevant strong interaction processes proceed and establish the hyperon population in the neutron star matter. Neutron star models are constructed at different levels of complexity starting from the most elementary one which assumes that neutrons are the only component. The more sophisticated version is formulated under the assumption that the neutron star matter  has to obey the constrains of charge neutrality and $\\beta$ equilibrium. Thus the model considered describes  high isospin asymmetric matter and it has to be extended by the inclusion of isovector-scalar  meson $a_{0}(980)$ ($\\delta $ meson)\\cite{delta}. For the sake of completness additional nonlinear vector meson interactions are included. When strange hadrons are taken into account uncertainties which are present in the description of nuclear matter are intensified due to the incompleteness of the available experimental data. The standard approach does not reproduce the strongly attractive hyperon-hyperon interaction seen in double $\\Lambda$ hypernuclei. In order to construct a proper model which do include hyperons the effects of hyperon-hyperon interactions have to be taken into account. These interactions are simulated via (hidden) strange meson exchange: scalar meson $f_0 (975)$ ($\\sigma^*$ meson) and vector meson $\\phi (1020)$ ($\\phi $ meson) and influence the form of the equation of state and neutron stars properties. \\newline The solution of the presented model is gained with the mean field approximation in which meson fields are replaced by their expectation values. The parameters used are adjusted in the limiting density range around saturation density $\\rho_0$ and in this density range give very good description in finite nuclei. However, incorporation of this theory to higher density require an extrapolation which in turn leads to some uncertainties and suffers of several shortcomings. The standard TM1 parameter set for high density range reveals an instability of neutron star matter which is connected with the appearance of negative nucleon effective mass due to the presence of hyperons. The Zimanyi-Moszkowski (ZM) \\cite{zima} model in which the Yukawa type interaction $g_{sN}\\varphi$ is replaced by the derivative one $(g_{sN}\\varphi/M_N)\\bar{\\psi}_N\\gamma_{\\nu}\\partial^{\\nu}\\psi$ exemplifies an alternative version of the Walecka model which improves the behaviour of the nucleon effective masses. It also influences the value of the incompressibility $K$ of neutron star matter. The derivative coupling effectively introduces the density dependence of the scalar and vector coupling constants. Knowing the form of the equation of state (EOS) is the decisive factor in determining properties of neutron stars such as: central density, mass-radius relation, crust extent or the moment of inertia.\\\\ The essential goal of this paper is to obtain within the described above model the equation of state for the neutron star matter on the basis of  calculations carried out for asymmetric nuclear matter in the relativistic mean field approach (RMF) and to compare the obtained results with ones that are relevant for high density calculations - namely with a quark star. The inclusion of $\\delta$ meson affects the neutron stars chemical composition changing the proton fraction which in turn affects the cooling mechanism. If the proton fraction is higher than the critical value of about $Y_p \\sim 0.11$ \\cite{lat} the direct URCA processes can proceed and this enhances the rate of neutron star cooling. Whether the proton fraction can exceed the critical value and at what density it occurs depends on the model. The proton fraction is almost entirely determined by the isospin-dependent part of the EOS thus the inclusion of $\\delta$ meson and nonlinear vector meson interactions  influence the neutron star structure and properties.\\newline This paper is organized as follows. Section 2 outlines the extended model with derivative coupling including hyperons and additional mesons, together with the collected equations of motions. Their solutions enable the construction of the equation of state . In section 3  the equilibrium conditions leading to the relevant hyperon star composition are presented together with the chosen values of parameters. Section 4 contains numerical results and the discussion of their influence on neutron stars properties. ",
        "conclusions": "In this paper the complete form of the equation of state of hyperon matter has been obtained with the use of the derivative coupling model in the framework of  an extended RMF theory which besides hyperons and leptons includes  the extended meson sector with additional  $\\delta$ meson and hidden strange mesons $\\sigma^*$ and $\\phi$. The model  considered is also supplemented with nonlinear vector meson interactions. This enlargement alters the symmetry properties of neutron star matter and through this neutron stars parameters. The value of baryon effective masses depend on the scalar meson condensates and at high densities when hyperon species appear the possibility of negative nucleon masses emerges. The derivative coupling model allows to avoid this difficulty reproducing reasonable value of  baryon effective masses for densities relevant for neutron stars. The inclusion of $\\delta$ meson and nonlinear vector meson interactions influences the chemical composition of a neutron star. This is especially evident  comparing the effective baryon masses in the density span $(3-4)\\times \\rho_0 $ and equilibrium compositions of the star. For the third group of parameters the higher value of asymmetry has been obtained. This changes the properties of a neutron star diminishing the hyperon core extent. The asymmetry of the system also influence the star radius, for the third group of parameters the one can obtain the lower value of the star radius. There is also possible a configuration representing a star with lower densities (a maximum radius configuration) which excluded the existence of a hyperon core. The model considered excludes a large hyperon fraction which can be connected with thermal properties of a hot neutron star. As the most populated strange baryon is the $\\Lambda$ hyperon it does not reduce significantly number of leptons in the star interior and thus the models calculated with the use of the parameter set I and II  do not exclude rapid cooling rate of the star. This is even more evident analyzing the proton fractions obtained for all parameter groups. The equilibrium proton fraction is also determined by the nuclear symmetry energy. The parameter sets I and II permit higher values of proton fraction which is indispensable for URCA processes to proceed. However, the equilibrium proton fraction $Y_p$ is significantly reduced for the third group of parameter.\\\\  Analyzing the gravitational binding energy one can come to the conclusion that  configurations with hyperons are energetically favorable than the one obtained with the use of TM1 parameter set for higher densities.\\\\  All assumption which have been made namely: the derivative coupling model  being connected  with the higher effective baryon masses, the inclusion of $\\delta$ meson and nonlinear vector meson interactions, and the repulsive nucleon-hyperon $\\Sigma$ interaction lead to the neutron star model with the value of maximum mass close to  $1.5 \\ M_{\\odot}$ with the reduced value of proton fraction and very compact hyperon core. The calculation of the quark matter equation of state. allows to construct the mass-radius relation for the quark star. Comparing gravitational binding energies one can come to the conclusion that the addition of hyperons to the model shifts the stable hyperon matter configuration towards higher densities even to the density range which is relevant for a quark star and at the same time makes the existence of a pure quark star more problematic."
    },
    "0212/astro-ph0212284_arXiv.txt": {
        "abstract": "We derive constraints on the mass-temperature relation of galaxy clusters from their observed luminosity-temperature relation and X-ray temperature function. Adopting the isothermal gas in hydrostatic equilibrium embedded in the universal density profile of dark matter halos, we compute the X-ray luminosity for clusters as a function of their hosting halo mass. We find that in order to reproduce the two observational statistics, the mass-temperature relation is fairly well constrained as $T_{\\rm gas} =(1.5\\sim 2.0)\\; {\\rm keV} (M_{\\rm vir}/10^{14}h_{70}^{-1}M_\\odot)^{0.5\\sim 0.55}$, and a simple self-similar evolution model ($T_{\\rm gas} \\propto M_{\\rm vir}^{2/3}$) is strongly disfavored. In the cosmological model that we assume (a $\\Lambda$CDM universe with $\\Omega_0=0.3$, $\\lambda_0=0.7$ and $h_{70}=1$), the derived mass-temperature relation suggests that the mass fluctuation amplitude $\\sigma_8$ is 0.7--0.8. ",
        "introduction": "While clusters of galaxies are relatively simple dynamical systems that consist of dark matter, stars, and X-ray--emitting hot gas, their thermal evolution is not yet fully understood. This is clearly illustrated by the well-known inconsistency of the observed X-ray luminosity-temperature ($\\Lx$-$T$) relation,  $\\Lx \\propto T^3$ \\citep[e.g., ][]{david93, markevitch98, ae99} against the simple self-similar prediction $\\Lx \\propto T^2$ \\citep{kaiser86}. Conventionally this is interpreted as evidence for preheating of intracluster gas;  additional heating tends to increase the temperature and the core size of the cluster, and to decrease the central density and the luminosity. Since the effect is stronger for less massive systems,  the slope of $\\Lx$-$T$ relation becomes steeper than that of the self-similar prediction \\citep{eh91, kaiser91}.  In addition, the mass-temperature ($M$-$T$) relation of clusters is also poorly determined. Although it is conventionally assumed that the gas shock heating is efficient enough and the temperature of the intracluster gas reaches the corresponding virial temperature of the hosting halos, this should be regarded as a simple working hypothesis. Nevertheless, the cosmological parameters derived from the cluster abundances are sensitive to the adopted $M$-$T$ relation. While this has already been recognized for some time (e.g., Figs.5d and 6d of Kitayama \\& Suto 1997), \\citet{seljak02} recently showed in a quantitative manner that the use of the observed $M$-$T$ relation by \\citet*{finogu} decreases the value of the mass fluctuation amplitude at $8\\; h^{-1}\\;$Mpc,  $\\sigma_8$, by $\\sim 20\\%$ where $h$ is the Hubble constant $H_0$ in units of 100 km s$^{-1}$ Mpc$^{-1}$ [we use, however, the dimensionless Hubble constant $h_{70} \\equiv H_{0}/(70\\;\\mathrm{km}\\; \\mathrm{s}^{-1}\\; \\mathrm{Mpc}^{-1})$ in the following analysis]. Therefore, the independent derivation of the cluster $M$-$T$ relation is important both in understanding the thermal history of the intracluster gas and in determining the cosmological parameters.  Our primary aim in this paper is to find the $M$-$T$ relation of clusters that reproduces the observed $\\Lx$-$T$ relation and X-ray temperature function (XTF). The reason we focus on $M$-$T$ relation is as follows:  since recent $N$-body simulations strongly indicate the universality of the density profile of the hosting halos of clusters (\\citealt*{nfw96}; \\citealt{moore98, js00}), the intracluster gas density profile in hydrostatic equilibrium with the underlying dark matter can be computed \\citep*{MSS98, SSM98} for a given mass of the halo. This enables one to make a reliable prediction for the $\\Lx$-$T$ relation once the $M$-$T$ relation is specified. In turn, one can obtain the $M$-$T$ relation that reproduces the observed $\\Lx$-$T$ relation without assuming an ad hoc model for the thermal evolution of intracluster gas.  In what follows, we parameterize the $M$-$T$ relation as a single power law, and derive the best-fit values of their amplitude and slope from the observed $\\Lx$-$T$ relation and the XTF. The result is compared with the recent observational studies by \\citet{finogu} and \\citet*{allen}. We also discuss the implications for the value of $\\sigma_8$ from cluster abundances. Throughout the paper, we adopt a conventional $\\Lambda$CDM model with density parameter $\\Omega_{0}=0.3$, cosmological constant $\\lambda_{0}=0.7$, dimensionless Hubble constant $h_{70}=1$,  and baryon density parameter $\\Omega_{{B}}=0.04\\;h_{70}^{-2}$. ",
        "conclusions": "We have presented the constraints on the (empirically parameterized) mass-temperature relation of galaxy clusters from the two different observational data, the luminosity-temperature relation and the temperature function.  We summarize in Figure~\\ref{fig:f08} our constraints on the $T_{\\rm gas, 0}$-$p_\\mathit{MT}$ plane, adopting the observed gas mass fraction (eq.~[\\ref{eq:fhotobs500}]). The fact that the simple self-similar evolution model $T_{\\rm gas} \\propto M_{\\rm vir}^{2/3}$ fails to explain the observed $\\Lx$-$T$ relation is well known and not at all new. Rather, it should be noted that the mass-temperature relation of galaxy clusters is fairly well constrained by combining the observed $\\Lx$-$T$ relation and XTF. It is encouraging that $T_{\\rm gas} =(1.5\\sim 2.0)\\; {\\rm keV} (M_{\\rm vir}/10^{14}\\;h_{70}^{-1}\\;M_\\odot)^{0.5\\sim 0.55}$ barely satisfy the two constraints simultaneously. This conclusion applies for $1\\le \\alpha \\le 3/2$ as long as the dark halo density profile is described by equation (\\ref{eq:nfw}).  In addition, our analysis implies that the amplitude of the mass variance $\\sigma_8$ in the standard $\\Lambda$CDM model should be 0.7--0.8. These values are significantly smaller than the previous estimates \\citep{vl96, eke96, KS97} but in better agreement with more recent results \\citep{seljak02, efsta}.   Since our current analysis has adopted a simple parameterized model for the mass-temperature relation, the result should be understood by a physical model of the thermal evolution of the intracluster gas. For instance, our result may be qualitatively explained by a kind of phenomenological heating of $T_{\\rm gas}(M_{\\rm vir}) = T_{\\rm vir}(M_{\\rm vir}) + 1\\; $keV. We plan to examine the implications for the possible heating sources from the derived mass-temperature relation of galaxy clusters using the Monte-Carlo modeling of merger trees."
    },
    "0212/astro-ph0212237_arXiv.txt": {
        "abstract": "{ We study the consequences  of antineutrino trapping in hot quark matter for quark star configurations with possible diquark condensation. Due to the conditions of charge neutrality and $\\beta$-equilibrium the flavor asymmetry increases with the number density of trapped antineutrinos. Above a critical value of the antineutrino chemical potential of $30$ MeV diquark condensation is inhibited at low densities and a two-phase structure emerges: a superconducting quark matter core surrounded by a shell of normal quark matter. When the quark star cools down below a temperature $T \\sim 1$ MeV, the mean free path of antineutrinos becomes larger than the thickness of the normal quark matter shell so that they get untrapped within a sudden process. By comparing the masses of configurations with the same baryon number we estimate that the release of energy due to the antineutrino untrapping transition can be in the range of $10^{51} \\div 10^{52}$ erg. ",
        "introduction": "\\label{sec:intro}  The engine which drives supernova explosions and gamma ray bursts being among the most energetic phenomena in the universe remains still puzzling [\\cite{Piran:2001da}]. The phase transition to a quark matter phase may be a mechanism that could release such an amount of energy [\\cite{Drago:1997tn}, \\cite{Berezhiani:2002ea}]. It has been proposed that due to the Cooper instability in dense Fermi gases cold dense quark matter shall be in the color superconducting state with a nonvanishing diquark condensate [\\cite{Alford:2000sx}, \\cite{Blaschke:2001uj}]. The consequences of diquark condensation for the cooling of compact stars due to changes in the transport properties and neutrino emissivities have been investigated much in detail, see [\\cite{Blaschke:1999qx}, \\cite{Page:2000wt}, \\cite{Blaschke:2000dy}, \\cite{Blaschke:2003yn}], and may even contribute to the explanation of the relative low temperature of the pulsar in the supernova remmant 3C58 [\\cite{cooling}].  Unlike the case of normal (electronic) superconductors, the pairing energy gap in quark matter is of the order of the Fermi energy so that diquark condensation gives considerable contributions to the equation of state (EoS) of the order of $({\\Delta}/{\\mu})^2$. Therefore, it has been suggested that there might be scenarios which identify the unknown source of the energy of $~10^{53}$ erg with a release of binding energy due to Cooper pairing of quarks in the core of a cooling protoneutron star [\\cite{Hong:2001gt}]. In that work the total diquark condensation energy released in a bounce of the core is estimated as $({\\Delta}/{\\mu})^2M_{{\\rm core}}$ corresponding to a few percent of a solar mass, that is $~10^{52}$ erg. In this estimate, general relativistic effects have been disregarded. It has been shown in [\\cite{Blaschke:2003yn}] by solving the selfconsistent problem of the star configurations, however, that these effects due to the stiffening of the EoS in the diquark condensation transition lead to an increase in the gravitational mass of the star  contrary to the naive estimates.  It has also been demonstrated [\\cite{Blaschke:2003yn}] that the energy release due to cooling of a quark core in a protoneutron star does not occur whithin an explosive process, since the diquark condensation is a second order phase transition.  In the present work, we propose a new  mechanism of energy release which involves a first order phase transition induced by antineutrino untrapping. (Anti-)neutrino trapping occurs in hot compact star configurations at temperatures $T\\geq 1$ MeV where the mean free path of (anti-)neutrinos is smaller than the typical size of a star [\\cite{Prakash:2001rx} and references therein].  During the collapse in the  hot era of protoneutron star evolution, antineutrinos are produced due to the $\\beta$-processes. Since they have a  small mean free path, they cannot escape and the asymmetry in the system is increased. This entails that the formation of the diquark condensate is shifted to higher densities or even inhibited depending on the fraction of trapped antineutrinos.   As the quark star cools, a two-phase structure will occur. Despite the asymmetry, the interior of the quark star (because of its large density) could consist of color superconducting quark matter, whereas in the more dilute outer shell, diquark condensation cannot occur and quark matter remains in the normal state, opaque to antineutrinos for $T\\geq 1$ MeV. When in the continued cooling process the antineutrino mean free path increases above the size of this normal matter shell, an outburst of neutrinos occurs and gives rise to an energy release of the order of $10^{51}-10^{52}$ erg. This untrapping transition is of first order and could lead to an explosive phenomenon.  The scenario to be detailed in the present paper suggests that  the first pulse of neutrinos emitted in the deleptonization stage of the core collapse, after a cooling time scale, shall be followed by a second pulse of antineutrinos as an observable characteristics of the present scenario. ",
        "conclusions": "We have investigated the effects of trapped antineutrinos on the asymmetry and  diquark condensates in a quark star configurations. By comparing configurations with fixed baryon number the release of energy in an antineutrino untrapping transition is estimated to be of the order of $10^{52}$ erg. Such a transition is of first order so that antineutrinos can be released in a sudden process (burst). This scenario could play an important r\\^ole to solve the problem of the engine of supernova explosions and gamma ray bursts. A second antineutrino pulse is suggested as an observable characteristics of the present scenario."
    },
    "0212/astro-ph0212371_arXiv.txt": {
        "abstract": "Within the framework of a model Universe with time variable space dimensions (TVSD), known as decrumpling or TVSD model, we study TVSD chaotic inflation and obtain dynamics of the inflaton, scale factor and spatial dimension. We also study the quantum fluctuations of the inflaton field and obtain the spectral index and its running in this model. Two classes of examples have been studied and comparisons made with the standard slow-roll formulae. We compare our results with the recent Wilkinson Microwave Anisotropy Probe (WMAP) data. ",
        "introduction": "One of the most intriguing challenges in modern physics is to find observable consequences of different kinds of theories in higher dimensions. We here present an inflationary model, known as decrumpling inflation model in which the number of spatial dimension has a dynamical behavior and decreases during the expansion of the Universe with a rate to be less than about $10^{-14} {\\rm yr}^{-1}$.\\cite{1}  Our motivation to study decrumpling inflation is to investigate cosmological implications of time variability of the number of spatial dimensions. To do so, we compute the spectral index and its running within the framework of decrumpling inflation. For more details about a model Universe with time variable space dimensions (TVSD), known as TVSD or decrumpling model, see Refs. [1-6].  To do this research an important conceptual issue is properly dealt with the meaning of time variability of the number of spatial dimensions. In Ref. [5], this conceptual issue has been discussed in detail.  Although time variability of spatial dimensions have not been firmly achieved in experiments and theories, such dynamical behavior of the spatial dimensions should not be ruled out in the context of cosmology and astroparticle physics.  Here, we will be concerened with the approaches proposed in the pioneer paper\\cite{12} where the cosmic expansion of the Universe is named decrumpling expansion and is due to decrease of the number of spatial dimensions.  The most important difference between decrumpling model and other attempts about the time evolution of spatial dimension is that in this model the number of extra spatial dimensions changes with time while in other theories the size of extra spatial dimensions is a dynamical parameter. Based on time variability of the size of spatial dimensions it has been reported\\cite{1} that the present rate of change of the mean radius of any additional spatial dimensions to be less than about $10^{-19} {\\rm yr^{-1}}$. It is worth mentioning that this result is based on dynamical behavior of the size of extra spatial dimensions while in decrumpling model we take the size of extra spatial dimensions to be constant and the number of spatial dimensions decreases continuously as the Universe expands. The present rate of time variation of the number of the spatial dimensions in decrumpling or TVSD model is about $10^{-13} {\\rm yr^{-1}}$.\\cite{1}  Another subject which lately has attracted much attention is the scalar spectral index and its running.  Primordial perturbations from inflation currently provide our only complete model for the generation of structure in the Universe. It is commonly stated that a generic prediction of inflationary models is a scale-invariant spectrum of adiabatic perturbations, characterized by a scalar spectral index $n_S$ that obeys $n_S-1=0$. However, this statement is only true for very special spacetimes like a pure de Sitter spacetime, which does not describe our cosmological history. For nearly all realistic inflationary models, the value of $n_S$ will vary with the wave number $k$.  Typically, since $n_S-1\\simeq 0$ on the scales probed by the cosmic microwave background (CMB), the deviations from a constant $n_S$ must be small. Nevertheless, increasingly accurate cosmological observations provide information about the scalar spectral index on scales below those accessible to anisotropy measurements of the CMB. Our such a wide range of scales, it is entirely possible that $n_S$ will exhibit significant running, a value that depends on the scale on which it is measured. Such running is quantified by the derivative $dn_S/d\\ln k$ and, in fact, recently released data from the WMAP satellite indicates that\\cite{7} \\begin{eqnarray} \\label{1} n_S\\;(k_0=0.05\\;\\;{\\rm Mpc}^{-1})&=&0.93^{+0.03}_{-0.03},\\\\ \\label{2} \\frac{dn_S}{d\\ln k}\\;(k_0=0.05\\;\\;{\\rm Mpc}^{-1})&=& -0.031^{+0.016}_{-0.018}. \\end{eqnarray}  The limits on the $n_S$ and its running using WMAP data alone are\\cite{8} \\begin{eqnarray} \\label{3} n_S\\;(k_0=0.002\\;\\;{\\rm Mpc}^{-1})&=&1.20^{+0.12}_{-0.11},\\\\ \\label{4} \\frac{dn_S}{d\\ln k}\\;(k_0=0.002\\;\\;{\\rm Mpc}^{-1})&=& -0.077^{+0.050}_{-0.052}, \\end{eqnarray} where $k_0$ is some pivot wave number.  Recent data, including that from the WMAP data satellite, show some evidence that the index runs (changes as a function of the scale $k$ at which it is measured) from $n_S>1$ (blue) on long scales to $n_S<1$ (red) on short scales. The authors of Ref.[9] investigated the extent to which inflationary models can accommodate such significant running of $n_S$. They show that within the slow-roll approximation, the fact that $n_S-1$ changes sign from blue to red forces the slope of the potential to reach a minimum at a similar field location.  Chaotic inflation models are usually based upon potentials of the form $V(\\phi)=a \\phi^b$ ($a$ is a constant and $b=2,4,...$ is an even integer). For the interesting cases of $b=2$ and $4$, it has been shown that\\cite{10} \\begin{eqnarray} \\label{5} \\frac{dn_S}{d\\ln k}=-0.8\\times 10^{-3}\\;(b=2,\\;{\\cal{N}}=50),\\\\ \\label{6} \\frac{dn_S}{d\\ln k}=-1.2\\times 10^{-3}\\;(b=4,\\;{\\cal{N}}=50), \\end{eqnarray} where ${\\cal{N}}$ is the e-folding number.  We will use the natural units system that sets $k_B$, $c$, and $\\hbar$ all equal to one, so that $\\ell_P=M_P^{-1}=\\sqrt{G}$. To read easily this article we also use the notation $D_t$ instead of $D(t)$ that means the space dimension $D$ is as a function of time.  The plan of this article is as follows. In section 2, we give a brief review of decrumpling or TVSD model. In section 3, we present the dynamical solutions of $\\phi(t)$, $a(t)$ and $D_t(t)$. In section 4, we first present the explicit and general formulae for the spectral index and its running within the framework of decrumpling or TVSD inflation and then apply them to two classes of examples of the inflaton potential. Finally, we discuss our results and conclude in section 5. ",
        "conclusions": "In previous our studies about time variable spatial dimension model\\cite{1} we showed that based on observational bounds on the present-day variation of Newton's constant, one would have to conclude that the spatial dimension of the Universe when the Universe was at the Planck length to be less than or equal to $3.09$. In [1] we concluded that if the dimension of space when the Universe was at the Planck scale is constrained to be fractional and very close to $3$, then the whole edifice of TVSD model loses credibility.\\cite{1}  In this paper, we have studied the effects of time variability of the spatial dimension on the time evolution of inflaton field, space dimension and scale factor for the potential $m^2\\phi^2/2$. We have also obtained general and explicit formulae for the scalar spectral index and its running in TVSD or decrumpling model and then applied them to two classes of examples of the inflaton potential. The correction terms due to time variability of space dimension depend on the e-folding number and the universal constant of TVSD model which is $C$. The numerical calculations for the spectral index and its running have been done for two classes of examples.  We have also discussed on dynamical solutions of inflaton field, scale factor and space dimension in TVSD or decrumpling chaotic inflation model. The outline of results has been shown in Figs. 1, 2 and 3. Fig.1 compares inflaton field within the framework of TVSD model and the constant three-space chaotic inflation. It is seen time variability of space dimension causes that inflaton field evolves slowly. This behavior of the scalar field can be seen by Eq. (\\ref{15}) where the effect of time variable dimension is like decreasing friction term in this equation, or subsequently causes that the scalar field changes slowly. In Fig. 2 the evolution of space dimension has been obtained. It is seen from Eq. (\\ref{11}) that temporal rate of space dimension is proportional to the square of space dimension. So during the inflationary epoch the Universe with $D_{P} = 10$ loses $6$ dimensions in comparison to $D_{P} = 4$ which loses less than one. In other words, $\\Delta D_{\\rm inflation}^{D_{P} = 10} = 6$ and $\\Delta D_{\\rm inflation}^{D_{P} = 4} <1 $. From Fig.3 it is seen that dynamical character of the space dimension increases the e-folding number.\\cite{4} Starting an Universe with higher dimensions at the Planck epoch, it leaves inflationary phase later.\\cite{3}  The final point that must be emphasized is about decrumpling model in cosmology. The original motivation of this model presented in the pioneer paper\\cite{12} was based on an {\\it ad hoc} assumption inspired from polymer physics. It is quite possible that this part of decrumpling model should be revised. However, just how this should be done is far from obvious. The progress in decrumpling model can only be made if there is a breakthrough in terms of finding a natural mechanism for varying the spatial dimension in some alternative fashion to that which we have considered."
    },
    "0212/astro-ph0212479_arXiv.txt": {
        "abstract": "The spectra of stars with the B[e] phenomenon are dominated by features that are related to physical conditions of circumstellar material around these objects and are not intrinsic to the stars. Because of this, they form a very heterogeneous group. This group contains objects with different evolutionary stages. Lamers et al. (1998) have suggested a new designation with five sub-groups, which indicate the evolutionary stage. They are: supergiants, pre-main sequence or Herbig Ae/Be, compact planetary nebulae, symbiotic and unclassified. The unclassified group has many objects that need a better study to resolve their evolutionary status. Forbidden lines can be a useful tool to solve this problem. They can give informations about chemical composition, ionization and density of the circunstellar medium and probably the evolutionary phase of these objects. We analize spectra of some galactic objects, obtained with FEROS and B\\&C spectrograph at 1.52 telescope in ESO (La Silla-Chile), with a special focus on the forbidden lines. We have studied the spectra of 5 B[e] stars of uncertain evolutionary stage. We find that one of them is a pre-WN star, the other four are supergiant B[e] stars. ",
        "introduction": "The presence of some forbidden lines is a criterion to distinguish differents groups of massive stars. For example, LBV and sgB[e] have similar spectral characteristics, however, [OI] lines are present only in the sgB[e] (Zickgraf, 1989). Below, we have a table showing the presence (Yes) or not (No) of some important forbidden lines. Figure 1 shows the profiles of forbidden and permitted NII lines in the spectra of HD 326823.  \\vskip 0.4truecm \\begin{center} {\\small \\begin{tabular}{|c|c|c|c|c|c|c|} \\hline \\multicolumn{1}{|c|}{ Objects } & \\multicolumn{1}{|c|}{ [OI] } & \\multicolumn{1}{|c|}{ [OII] } & \\multicolumn{1}{|c|}{ [OIII] } & \\multicolumn{1}{|c|}{ [SII] } & \\multicolumn{1}{|c|}{ [NII] } & \\multicolumn{1}{|c|}{ [FeII] } \\\\ \\hline \\hline HD 87643 & Yes & No & No & Yes & No & Yes\\\\ Hen 3-847 & Yes & No & No & Yes & Yes & Yes \\\\ GG Car & Yes & No &  No & No & Yes & Yes \\\\ MWC 300 & Yes & No & No  & Yes  & Yes & Yes \\\\ HD 326823 &  No & No & No & Yes & Yes & Yes \\\\ \\cline{1-3} \\hline \\end{tabular} } \\end{center}  \\vskip 60truecm \\vbox{\\special{psfile=fig1.ps angle=0 vscale=27 hscale=37 voffset=-190 hoffset=65}} \\vskip 5.5truecm \\noindent {\\footnotesize {\\bf \\it Figure 1 -}Profiles of NII lines in the high resolution spectrum (FEROS-agreement ESO/ON) of HD 326823. The permitted lines are in absorption, indicating a photospheric origin, and the forbidden are in emission, indicating a circumstellar origin. } ",
        "conclusions": ""
    },
    "0212/astro-ph0212153_arXiv.txt": {
        "abstract": " ",
        "introduction": "Extremely Red Objects (EROs), which have optical-to-infrared colors which differ significantly from typical field sources, encompass a wide variety of phenomen\\ae.  Galaxies of assorted types make up the dominant component of ERO samples, but one can also find low mass stars, gravitationally lensed sources, and transient sources such as variable stars, asteroids or supernovae which may not be initially recognized as such.  The term Extremely Red Galaxies (ERGs) is also commonly used, sometimes interchangeably, but usually refers to a sample of EROs which has been cleaned of the objects which are not galaxies.  ERGs are therefore a subset of the EROs.  We adopt the more general term ERO throughout this paper.  Both the definition and interpretation of EROs has evolved somewhat since their initial discovery, and it is useful to review the subject here for some historical perspective.  When first identified as a distinct population of sources ~\\citep{ERR88}, EROs were thought to be good candidates for primeval galaxies.  Subsequent observations ~\\citep{ERR89}, however, showed these early EROs to be $z \\sim 0.8$ elliptical galaxies.  Additional EROs were noted in the following years \\citep{MPW92,ED92,Pea93,Graham94,HR94,Sea94,DSD95,Dj95,Treu98,Im02}.  Most of these were serendipitous detections, identified on images targeting known, high-redshift radio galaxies or other active galactic nuclei.  Little or no followup work was done on these objects at the time, which to some extent reflected the limited capabilities of existing instruments and telescopes.  These EROs were identified with colors spanning a wide range: ($R-K$)\\,$\\geq$\\,5--7, or ($I-K$)\\,$\\geq$\\,4--6.  A resurging interest in EROs accompanied the development of the Submillimeter Common User Bolometer Array \\citep[SCUBA]{SCUBA}, and the subsequent detection of the extremely red galaxy HR\\,10 ~\\citep[this source is also known as HR94\\,10 or ERO\\,J164502+4626.4]{HR94} at 850\\,$\\mu$m ~\\citep{cim98,dey99}.  At a redshift of 1.44 ~\\citep{GD96}, the detection of HR\\,10 in the submillimeter implied the presence of massive quantities of dust accompanied by very high star formation rates. The ERO population was thought to provide fertile hunting grounds for more submillimeter-bright galaxies at high redshift. Additional observations have not supported this idea, however, with only a relatively small fraction, on the order of 20\\% ~\\citep[Thompson, priv. comm.]{and99,mohan02}, of the bright EROs ($K < 19.5$) showing strong submillimeter emission.  The development of larger-format infrared arrays and wider field instrumentation enabled subsequent field surveys to cover enough area to assemble significant samples of systematically selected EROs for further study, in blank fields ~\\citep{tho99,daddi00a,LCIR01} as well as targeted surveys ~\\citep{Chap00,Cim00,Liu00}.  \\citet{tho99} adopted a color selection for EROs of ($R-K$)\\,$>$\\,6\\fm0.  The motivation was that this color was redder than the expected colors of elliptical galaxies with anything but the highest formation redshifts ($z_f > 10$), and thus represents an extreme color for any normal galaxy. The assumption at the time was that the extremely red galaxy population consisted of either old ellipticals or young, dusty starbursts ~\\citep{cim98,tho99,dey99}.  The relative contribution of these two types of galaxies would have a bearing on the timing of massive galaxy formation and their subsequent evolution.  It is important to emphasize that, at that time, the term ``young, dusty starbursts'' referred specifically to massive starbursts like that seen in HR\\,10 or luminous infrared galaxies. Multi-band photometry could potentially distinguish between ellipticals and starbursts ~\\citep{PM00}, but this technique requires very low photometric uncertainties to work well (see, for example, ~\\citet{Mea02}).  In order to better study the $z \\sim 1$ elliptical galaxy population, \\citet{daddi00a} adopted a bluer color selection limit, ($R-K$)\\,$>$\\,5.3, set by the expected colors of a $z = 1$ passively evolving old stellar population.  This definition, or the roughly equivalent ($I-K$)\\,$>$\\,4.0, for the ERO color selection criterion has generally been adopted in the majority of subsequent work.  While there are a number of redshifts now known for EROs ~\\citep{GD96,Sea99,Liu00,afonso01,smith02b}, systematic redshift surveys of complete samples are only now becoming available ~\\citep{cim02}.  Morphological information based on high-resolution HST imaging for complete samples of EROs are also only now starting to appear ~\\citep[this work]{smith02a}.  Without similar spectroscopic or morphological information, earlier ERO surveys divided the ERO population into two components: old, evolved systems or dusty, massive starbursts.  But the true nature of K-selected EROs is likely to be much more complex, as suggested by recent work ~\\citep{LCIR01,cim02}.  ~\\citet{LCIR01} find a large scatter in the $(V - I)$ colors of their ERO sample, best fit by passive evolution models with extended star formation ($\\tau = 1$~Gyr).  This implies that the star formation history of EROs is more complex than a binary division into evolved ellipticals or dusty, massive starbursts implies.  From their K20 survey, ~\\citet{cim02} found that about half of their spectroscopic sample of $\\sim$30 EROs are dusty star-forming galaxies with emission lines, while the remaining half are old stellar populations with absorption line spectra.  However, the simple presence of line emission could span a wide range of galaxy types, from bulge-dominated, late type spiral galaxies with a small amount of star formation through the more massive starbursts like HR\\,10.  Dust could also completely obscure any on-going star formation, to the point that the optical and UV emission lines are not seen ~\\citep{PW00}.  Examples of what appear to be quiescent disks at $z \\sim 1.5$ exist ~\\citep{dokk01,smith02b}.  There are important differences in the formation and evolution of quiescent normal galaxies and young, dusty, massive starbursts. Morphologies have the potential to distinguish between the various interpretations, which motivated this work.  In this paper, we present the high resolution morphologies derived from HST WFPC2 images for a large sample of K-selected EROs. Our results reveal for the first time a new type of ERO which dominates the population and is neither an early-type galaxy nor a dusty, massive star forming galaxy.  We will also discuss the implications of our results for the past and future evolution of massive galaxies at $z \\sim 1$. ",
        "conclusions": "\\subsection{Morphological Distribution}  Our morphologies are based on F814W images.  Assuming the median redshift of one for EROs from the ~\\citet{cim02} sample with K$ \\leq 19.2$, the WFPC2 data sample a rest-frame wavelength of 4100\\,\\AA.  The F814W images thus represent a compromise between sensitivity to star formation at shorter restframe wavelengths and better probing any extended old stellar populations at longer restframe wavelengths.  As shown in simulations by ~\\citet{hibvac97}, morphological classifications using F814W images do not show any significant biases at $z \\sim 1$.   Using the results from our visual classification, we find that 30$\\pm$5\\% of our EROs have morphologies consistent with spheroidal (B and BD) galaxies. Disks (D and DB) dominate the EROs at 64$\\pm$7\\%\\ of the sample.  Only 6\\% of the EROs were unclassifiable, due primarily to low SNR on the WFPC images.  The uncertainties are derived simply from the square root of the number of EROs in that subset, and does not try to include the unclassifiable sources.  We plot in Figure~\\ref{relfrac} the relative fractions of spheroids and disks in our visual classification as well as from the MDS profile fitting.  \\begin{figure}[!ht] \\plotone{f7.eps} \\caption{Relative fraction of spheroidal galaxies vs. disks as classified in this survey (gray bars) and by the MDS automated profile fitting (white bars).  Not included in this plot are the 6\\% of sources unclassifiable by us or the 14\\% unclassified by the MDS.  \\label{relfrac}} \\end{figure}  We find that the relative morphological mix of EROs in subsets constructed of the EROs in known foreground galaxy cluster fields vs. those selected in the remaining ``blank'' fields is consistent to within the uncertainties. Since any lensing should be unbiased with respect to background galaxy morphologies, we do not differentiate between cluster and field sources in the remaining analysis.  Some of the EROs, including both disks and spheroids, appear to be involved in recent mergers or show evidence of strong interactions (e.g. tidal tails, strong asymmetries).  These comprise 17$\\pm$4\\% of the sample.  Without additional information, these systems represent the most likely source of possible large-scale starbursts. This is consistent with the observation that among {\\it bright} EROs, roughly 20\\% have 850$\\mu$m detections ~\\citep{and99}. One third of the galaxies we classed as spheroids have one or more faint companions or show signs of recent interaction, suggesting that a significant fraction of otherwise old stellar populations may not be in purely passively evolving systems.  The morphological classifications derived from the MDS profile fitting are largely consistent with our visual results, but show a slightly higher fraction of spheroidal galaxies, with 37$\\pm$6\\% of the sample classed as B or BD (with a bulge fraction larger than 50\\%).  There is a corresponding drop in the disk fraction, with 50$\\pm$7\\% classed as DB or D, but disks still dominate the overall $\\rm (I-K)\\geq 4$ ERO population.  A larger fraction of the sample, 14$\\pm$3\\%, were not fit with the bulge$+$disk models due to low SNR or a more conservative avoidance of the CCD edges in the WFPC2 data.  We find that our ERO sample selected with $(F814W - Ks) \\geq 4$ is dominated by disk galaxies, and not by spheroids or strongly interacting systems. The star formation rates in these disks could span a wide range, including normal galaxies with fairly quiescent star formation.  HR\\,10 type systems, with mostly young stars and undergoing massive starbursts, may comprise only a small fraction of the sample.  The origins of their red colors may be traced to a significant old stellar population combined with some dust extinction, especially considering the edge-on orientation of many of the disks in our sample of EROs.  Even though most of the stellar mass may already be in place by $ z \\sim 1 - 2$ for these EROs, such a large fraction of disk galaxies implies that there is still a substantial amount of gas available to feed on-going star formation.  The existence of such a large fraction of disk galaxies and interacting systems in our sample suggests that hierarchical merging may be an important mechanism for the formation and evolution of the ERO population. However, the scenario in which ellipticals were formed in a ``monolithic collapse'' at high redshifts and evolve passively thereafter cannot be excluded.  While 30\\% of our ERO sample are clearly early type galaxies, whose colors are consistent with old stellar populations formed at high redshifts. We also found that one-third of the spheroids in our sample have faint companions or signs of interaction. This suggests that although a majority of their mass could be assembled rapidly at high redshift, these systems are not simply isolated, passively evolving old stellar populations, and continuing accretion of gas or merger events plays a significant role in their evolution. Examples of secondary star formation in field E/S0s ($z \\sim 0.1 - 0.73$) can be found in ~\\citet{treu02}.  Our results contradict those of ~\\citet{mor00}, which are also based on HST morphologies.  They find that 50--80\\%\\ of their sample have E/S0 morphologies on the basis of one-component exponential model fits.  However, we note that their ERO sample was assembled from the published literature, with the corresponding heterogeneous selection functions of the original surveys.  In addition, the morphologies were determined on HST images from both WFPC2 and NICMOS, probing widely different rest-frame wavelengths and thus differing sensitivities to star formation or old stellar population.  \\subsection{Surface Density} \\label{SDsect}  The integrated surface density of EROs, is defined as total number of EROs brighter than a given magnitude per unit area on the sky. Because the area covered in this survey is a function of the magnitude, the differential number of EROs selected in each magnitude bin is a function of both the magnitude and the area surveyed.  Calculating the integrated surface density thus required rescaling the number of EROs in each of the brighter bins by the appropriate area ratio prior to integration.  We plot the resulting surface density of EROs from our survey derived by this method in Figure~\\ref{surfdens} (filled diamonds). The uncertainties were derived simply from the square root of the rescaled number of EROs.  \\begin{figure}[!ht] \\plotone{f8.eps} \\caption{Cumulative ERO surface density from this work (filled diamonds), as well as several other surveys for comparison.  The uncertainties are derived assuming only Poisson counting statistics in each bin. \\label{surfdens}} \\end{figure}  For comparison, we also plot in Figure~\\ref{surfdens} the results from other recent surveys ~\\citep{tho99,smith02a,daddi00a,LCIR01,cim02,barg99}. The primary difficulty in making such comparisons is that each survey used a different set of filters and different selection criteria for identifying EROs.  In order to make a general comparison between these disparate surveys, we make several simplifying assumptions.  First, we treat all K filters as functionally equivalent (e.g. $K == K^\\prime == K_s$), so no color terms are applied to convert magnitudes between them.  The same is true for the several different I filters and R filters used.  We note that these assumptions are generally made in most recent published ERO surveys unless the color selection limit is explicitly tied to some fiducial model SED, typically a passively evolving $z=1$ old stellar population (i.e. elliptical galaxies), {\\em and} the specific filters used for the survey.  Second, we adopt the generic colors of ($R-I$)$=$1\\fm3 and ($H-K$)$=$1\\fm0 to convert the surveys based on R band or H band data to our $I-K$ colors.  Third, we plot without additional correction the ~\\citet{cim02} results, even though they select EROs at a bluer limit: ($R-K_s$)$\\geq$5\\fm0.  Finally, we adopt the ~\\citet{barg99} results directly for ($I_c-HK^\\prime$)$\\geq$3\\fm7, which they show to be equivalent to ($I-K$)$\\geq$4\\fm0 for that filter set.  While the generic color conversions we adopt obviate the possibility of more detailed comparisons between the different ERO surveys, they do allow us to consider broad trends in the surface densities as a function of the selection filters.  First, the two I-band based surveys (this work and ~\\citet{LCIR01}) agree quite well over the region $17^{\\rm m} \\leq K_s \\leq 19^{\\rm m}$.  This suggests that our sample is neither significantly inhomogeneous nor incomplete over this range, although the turnover in counts to fainter magnitudes suggests incompleteness in our sample for $K>19^{\\rm m}$.  The two larger R-band based surveys ~\\citep{tho99,daddi00a} also agree with each other over this same range in K-band magnitude, but are systematically about a factor of three lower than the I-band surveys.  The lensing-corrected surface density of EROs from the cluster-pointed survey of ~\\citet{smith02a}, despite using an R-band filter, agrees well with the I-band based surveys at brighter magnitudes ($K<19^{\\rm m}$). The other two surveys, ~\\citet{cim02,barg99}, differ from the other results.  Aside from the uncertainties in converting from one filter system to another, there are two primary effects which can qualitatively account for the similarities and differences in the surface densities of EROs from these different surveys: cosmic variance, and color selection effects.  Both the ~\\citet{cim02} and ~\\citet{barg99} surveys cover relatively small, connected areas on the sky, and are thus more subject to cosmic variance, especially considering the strong clustering seen in ~\\citet{daddi00a}.  The ~\\citet{smith02a} survey is composed of 10 widely-separated sight lines and thus should be less sensitive to cosmic variance, but their sample does show a wide field-to-field variation in the number of EROs.  ~\\citet{cim02} selects EROs at a bluer limit, ($R-K_s$)$\\geq$5\\fm0, which likely contributes to their higher ERO surface density.  A color selection effect may contribute to the apparent differences in surface density between R-band based ~\\citep{tho99,daddi00a} and I-band based (this work and ~\\citet{LCIR01}) ERO surveys.  We offer below (see \\S\\ref{ColSel}) a qualitative argument on this, as we do not have a proper multiband deep survey to address this with real data.  \\subsection{Volume Density}  We can make an estimate of the volume density of EROs with some assumptions on the range of redshifts at which they may be found, and compare these results with the local density of massive galaxies. The color selection limit of ($I-K$)\\,$\\geq$\\,4\\fm0 sets the lower redshift bound to $z = 1$, which is appropriate for passively evolving old stellar populations.  However, photometric uncertainties could make this $z = 1$ boundary somewhat fuzzy.  Galaxies with significant dust extinction (see the following section) could also lie at lower redshifts. \\citet{cim02} obtained spectroscopic redshifts for a sample of EROs with ($R-K$)$>$5 which includes sources of both types down to $z \\sim 0.7$. We adopt an upper cutoff to the assumed redshift range of $z = 2$, as higher redshift EROs would be anomalously luminous given their bright $K$-band magnitudes.  Under the above assumptions, we derive a co-moving volume covered by our survey to be $4\\times10^5\\,h_{70}^{-3}$\\,Mpc$^3$.  This volume is only weakly dependent on the assumed redshift range, and only changes by a factor of two if the redshift range is narrowed to $z \\sim 1.0 - 1.5$ or broadened to $z \\sim 0.7 - 2.8$.  This volume was derived from the total survey area of 409 square arcminutes, and does not take into account the variable survey depth with magnitude.  The galaxies classed as spheroids in our survey have a co-moving volume density of $1\\times10^{-4}\\,h_{70}^3$\\,Mpc$^{-3}$. The disks have a co-moving volume density about twice as large, and the total ERO sample (115 EROs) reach a density of $3\\times10^{-4}\\,h_{70}^3$\\,Mpc$^{-3}$.  To compare with nearby massive galaxies, we adopt the local K-band luminosity functions for early-type and late type galaxies from ~\\citet{Kochanek01}.  We integrate from 10L$^*$ down to 1L$^*$, which corresponds to our K-band limit at $z = 1$ after correcting for passive evolution and cosmological K-corrections.  We find that the EROs can account for only one-third of the local massive galaxies, and that the relative morphological mix is about the same in the two samples.  This is reasonable, but should be considered only an order-of-magnitude agreement given the factors of $\\sim$2 uncertainties in the volume densities arising from the assumptions on the redshift distributions, area surveyed as a function of depth (\\S\\ref{sec:area}), and contamination (\\S\\ref{sec:eodisks}).  \\subsection{Dust Extinction in Disky EROs} \\label{sec:eodisks}  While classifying the EROs, we also noted that 40\\% of the disky systems (DB and D) appeared to be sufficiently edge-on that even small amounts of dust in a disk could have a disproportionately large effect on the overall system color.  These systems are noted in Table~\\ref{ERO_tab} with italicized entries under morphology: {\\em DB} and {\\em D}, and we show two examples of edge-on EROs in Figure~\\ref{edgeon}.  This is far more than expected from a set of randomly oriented galaxies, suggesting that orientation effects are responsible for their inclusion in the ERO sample.  \\begin{figure}[!ht] \\plotone{f9.eps} \\caption{Examples of two edge-on disks where dust extinction may contribute significantly to their overall ($I-K_s$) color.  Each image is 8'' square. The labels correspond to their source numbers from Table~\\ref{ERO_tab}. \\label{edgeon}} \\end{figure}  Given the potentially large extinctions possible from dust in otherwise normal disk galaxies, it is possible that the edge-on systems are at lower redshift ($z < 1$).  Several of these show extended disks of large apparent size (several arcseconds), which would be unusually large (tens of kiloparsecs) if at $z \\sim 1$ or more.  The edge-on systems comprise half of the disky EROs, or one-third of the total ($I-K$) selected ERO sample.  They thus represent a large and previously unanticipated source of contamination in the ERO population.  \\subsection{Color Selection Effects} \\label{ColSel}  How comparable are ERO samples selected using an $(R - K) \\ge 5.3$ color versus an $(I - K) \\ge 4$ color?  This important issue has never been clearly addressed.  We investigated this issue using model spectral energy distributions, but lack the necessary multiband data to compare to the models.  Deep, wide-field infrared/optical surveys should be able to address this point in more detail.  In Figure~\\ref{modelRIK} we plot the ($R-K$) vs. ($I-K$) colors for a ~\\citet{BC96} model approximating an old stellar population (OSP, $\\tau = 0.1, z_f=30$) or passively evolving elliptical galaxy.  We also plot two models with longer exponential decay times ($\\tau = 1.0$) but differing formation redshifts ($z_f = 30, 5$), which should contain a significant fraction of old stars in the range of $1 < z < 2$, but still have some residual star formation.  A similar plot covering the $VIH$ color-color plane can be found in ~\\citep[their Figure 2]{LCIR01}. Their data show that EROs selected with ($I - H$)\\,$\\geq$\\,3 have a wide scatter in the $(V - I)$ color, which the authors interpreted as due to prolonged star formation.  \\begin{figure}[!th] \\plotone{f10.eps} \\caption{$RIK$ colors for a ~\\citet{BC96} model approximating an old stellar population (OSP, $\\tau = 0.1, z_f=30$), as well as two models with longer exponential decay times ($\\tau = 1.0$) but differing formation redshifts ($z_f = 30, 5$).  All three curves are marked with squares at $z \\sim 1$.  Note that that pure, passively-evolving OSPs would be selected in either an ($R-K$) or ($I-K$) survey for EROs, but that any residual star formation, such as from a disk, would quickly drop objects out of the ($R-K$) sample.  Reducing the redshift of formation much below five, or increasing the residual star formation rates ($\\tau > 1$) would drop objects out of both samples.  Dust counteracts these effects somewhat, shifting objects towards the upper right of this plot. \\label{modelRIK}} \\end{figure}  The two dotted lines in Figure~\\ref{modelRIK} mark the fiducial colors ($R - K$)\\,$=$\\,5\\fm3 and ($I - K$)\\,$=$\\,4\\fm0, representing the colors of a $z = 1$ passively-evolving elliptical galaxy used by most surveys to select EROs.  The squares mark the $z \\sim 1$ points in each model curve.  Assumptions on the models used, as well as the assumed cosmology and the specific filter bandpasses used for ERO surveys can account for several tenths of a magnitude variation in the expected colors of a $z = 1$ passively-evolving elliptical galaxy.  Several predictions can be made from this color selection effect.  First, is that both R-band and I-band based ERO surveys should select the {\\em same population} of passively-evolving old stellar populations (elliptical galaxies).  Second, I-band based ERO surveys should preferentially include disk galaxies.  Light from a bulge comprised of older stars would dominate the ($I-K$) color, while even small amounts of residual star formation in the disk keeps them too blue in ($R-I$) to be included in an ($R-K$) selected sample.  Other factors, such as dust extinction, may counteract the star-formation and contribute significantly to their overall color.  This implies that ERO samples selected on their ($R - K$) color may not be comparable to samples selected on their ($I - K$) color.  Figure~\\ref{modelRIK} can also be used to set some constraints on the formation redshift for the EROs.  Models with exponential decay times longer than about 1.5\\,Gyr ($\\tau \\geq 1.5$), or with a formation redshift lower than five ($z \\leq 5$) simply have too much residual star formation.  Without any reddening from dust, the blue light from a young population of stars would be sufficient to drop these galaxies out of either  ($R-K$) or ($I-K$) samples of EROs.  Thus, the brighter EROs classed as spheroids, especially those without any evidence of a disk or ongoing star formation, are likely to have formed a majority of their stars at relatively high redshift ($z > 5$).  The morphological mix in our sample of relatively bright, ($I-K$)-selected EROs is similar to that of ~\\citet{smith02a} among a sample of fainter EROs ($K \\leq 20.6$, $R-K \\geq 5.3$) identified in the fields of foreground clusters massive enough to gravitationally lens the higher redshift EROs. They classify 18\\% of their sample as compact, and 50\\% as irregulars (including disk-like systems), while 32\\% are too faint to be classified.  Considering redder subsamples from both this work ($F814W - K_s \\geq 5$) and ~\\citet{smith02a} ($R - K \\geq 6.0$), we again find similar results. Of the 11 redder EROs in our sample, nine were classifiable.  Of these nine, 89$\\pm$31\\% have disk morphologies, a much higher fraction than in the full sample.  The redder ~\\citet{smith02a} subsample has $\\sim$90\\% with disk/irregular morphologies.  The expectation from the color selection effect is that (R-K) samples should contain a higher fraction of spheroids at a given $K$ magnitude limit.  However, the fainter ~\\citet{smith02a} sample, about three magnitudes deeper than our ERO sample, is composed primarily of irregulars and disks.  This suggests that the morphological mix of EROs does change at fainter magnitudes, and ~\\citet{smith02a} conclusion that fainter and redder samples are dominated by massive, dusty starbursts does not contradict our findings.  Clearly, this color selection effect needs to be investigated further, with larger and deeper samples of EROs with both (R-K) and (I-K) selection on the same area of sky, so that a proper comparison can be made."
    },
    "0212/astro-ph0212365_arXiv.txt": {
        "abstract": "The wealth of high quality data now available on the M87 jet inspired us to carry out a detailed analysis of the plasma physical conditions in the jet.  In a companion paper (Lobanov, Hardee \\& Eilek, this proceedings) we identify a double-helix structure within the jet, and apply Kelvin-Helmholtz stability analysis to determine the physical state of the jet plasma. In this paper we treat the jet as a test case for {\\it in situ} particle acceleration.  We find that plasma turbulence is likely to exist at levels which can maintain the energy of electrons radiating in the radio to optical range, consistent with the broadband spectrum of the jet. ",
        "introduction": "We know a great deal about the jet in M87.  Radio \\citep{OHC} and optical \\citep{HST} images reveal ordered, filamentary structures whose synchrotron emission extends from radio at least to optical frequencies, and probably up to X-rays \\citep{Xray}.  Multi-epoch studies detect relativistic proper motion, at $\\gamma$ of a few\\citep{Biretta}.  The jet begins by expanding uniformly, out to $\\sim 2$ kpc (we assume an angle $\\sim 40^{\\circ}$ to the line of sight, and a distance of 17 Mpc). At that point, the location of the bright knot A, it recollimates, and continues for another 1-2 kpc before disrupting strongly and depositing its matter and energy in the larger radio halo. The minimum pressure required to produce the synchrotron emission remains approximately constant during the expansion. Because $p_{min}$ measures the energy density in relativistic electrons and magnetic field, $p_{min} \\propto u_e^{4/7} u_B^{3/7}$, we know that these quantities do not decay adiabatically during the expansion.  This is clear evidence that {\\it in situ} energization is occuring.  The broadband, slice-integrated synchrotron spectrum is nearly constant along the jet. This suggests that the electron energy distribution  is also nearly constant along the jet, and  supports the idea that it is maintained against losses by {\\it in situ} energization. We must recall, however, that the jet is not internally homogeneous. The radio-bright filaments probably show us high-field regions. In addition, resolved two-dimensional images show that the optical knots are more concentrated than the radio knots.  We have enough information about this jet to justify detailed modelling of the turbulence and its effect on the particles. To be specific, we consider particle acceleration by Alfvenic turbulence. In this paper we summarize our analysis, which will be published elsewhere in more detailed form. ",
        "conclusions": ""
    },
    "0212/astro-ph0212492.txt": {
        "abstract": "Improving our understanding of the initial conditions and earliest stages of protostellar collapse is crucial to gain insight into the origin of stellar masses, multiple systems, and protoplanetary disks. Observationally, there are two complementary approaches to this problem: (1) studying the structure and kinematics of prestellar cores observed prior to protostar formation, and (2) studying the structure of young (e.g. Class 0) accreting protostars observed soon after point mass formation. We discuss recent advances made in this area thanks to (sub)millimeter mapping observations with large single-dish telescopes and interferometers. %Two results have emer-ged. First, In particular, we argue that the beginning of protostellar collapse is much more violent in cluster-forming clouds than in regions of distributed star formation. %Second, protostars in clusters seem to originate from %finite, detached reservoirs of mass, suggesting the IMF is largely %determined by pre-collapse cloud fragmentation. Major breakthroughs are expected in this field from future large submillimeter instruments such as Herschel and ALMA. ",
        "introduction": "Although the formation of low-mass stars is now reasonably well understood in outline (see, e.g., Mannings, Boss, \\& Russell 2000 for recent reviews), several important aspects remain poorly known, such as the initial stages of the collapse process, the mechanism(s) selecting stellar masses, or the formation of multiple systems. Some progress on the earliest stages of star formation was achieved over the last decade thanks to the use of sensitive receivers on large (sub)millimeter radiotelescopes such as JCMT, CSO, and the IRAM 30m. Young protostars were identified at the beginning of the main accretion phase (Class 0 objects -- Andr\\'e, Ward-Thompson, \\& Barsony 1993), and the starless dense cores extensively studied in NH$_3$ by Myers and collaborators (e.g. Benson \\& Myers 1989) were found to be characterized by flat inner density gradients (Ward-Thompson et al. 1994). Direct evidence for infall motions was observed toward a large number of Class~0 protostars and prestellar cores (e.g. Gregersen et al. 1997 -- see \\S ~3 below and Evans, this volume). Class~0 objects were also found to drive more powerful outflows than more evolved (Class~I) protostars, suggesting a marked decrease of the mass accretion/ejection rates in the course of protostellar evolution (Bontemps et al. 1996). As advocated by Henriksen, Andr\\'e, \\& Bontemps (1997), there may be a causal relationship between these results, in the sense that the accretion/ejection decline during the protostellar phase may be a direct consequence of the form of the density profile at the prestellar stage.\\\\ Further studying the detailed properties of prestellar cores and young protostars is of prime importance to distinguish between collapse models and shed light on the origin of stellar masses. Indeed, the effective reservoirs of mass available for the formation of individual stars may be largely determined at the prestellar stage, and it is during the main protostellar accretion phase that stars accrete some fraction of these reservoirs and build up the masses they will have on the zero-age main sequence.  \\begin{figure}[ht] %\\resizebox{\\hsize}{!}{\\includegraphics{/home/storage/pandre/tmp3/fig/l1544_iram%04191.ps}} \\resizebox{\\hsize}{!}{\\includegraphics{andrep1_rev.ps}} \\caption{Dust continuum maps of L1544 (a) and IRAM~04191 (b) at 1.3~mm taken with the IRAM 30~m telescope and the MPIfR bolometer array (from Ward-Thompson et al. 1999 and Andr\\'e et al. 1999, respectively). Effective resolution: 13\\arcsec ; base contour and contour step: 20~mJy/beam. The direction of the magnetic field measured in L1544 with the SCUBA polarimeter on JCMT (Ward-Thompson et al. 2000) and the collimated CO(2-1) bipolar flow emanating from IRAM~04191 are shown. \\label{fig-intro}} \\end{figure}  After a brief introduction of the theoretical background (\\S ~1.1), we review recent observational advances concerning the density and velocity structure of cloud cores in \\S ~2 and \\S ~3, respectively. We conclude in \\S ~4 with a comparison between observations and theoretical models. %\\vskip 0.3cm  \\subsection{Collapse Initial Conditions: Theory}  The inside-out collapse model of Shu (1977), starting from a singular isothermal sphere (SIS) or toroid (cf. Li \\& Shu 1996, 1997), is well known and underlies the `standard' picture of isolated, low-mass star formation (e.g. Shu, Adams, \\& Lizano 1987).\\\\ Other collapse models exist, however, which adopt different initial conditions. In particular, Whitworth \\& Summers (1985) have shown that there is a two-parameter continuum of similarity solutions to the problem of isothermal spherical collapse. One of the parameters measures how close to hydrostatic equilibrium the system is initially, while the other parameter reflects how important external compression is in initiating the collapse. In this continuum, the solutions proposed by Shu (1977) and Larson (1969)-Penston (1969) represent two extreme limits. All of the similarity solutions share a universal evolutionary pattern.  At early times ($t < 0$), a compression wave (initiated by, e.g., an external disturbance) propagates inward %at the sound speed, leaving behind it a $\\rho(r) \\propto r^{-2}$ density profile. At $t = 0$, the compression wave reaches the center and a point mass forms which subsequently grows by accretion.  At later times ($t > 0$), this wave is reflected into a rarefaction or expansion wave, propagating outward through the infalling gas (at the isothermal sound speed $a_s$), and leaving behind it a free-fall $\\rho(r) \\propto r^{-1.5}$ density distribution. Several well-known features of the Shu model (such as the expansion wave) are thus in fact common to all solutions. The various solutions can be distinguished by the {\\it absolute} values of the density and velocity at $t \\sim 0$. In particular, the Shu (1977) solution has $\\rho(r) = (\\as^2/2\\pi\\,G)\\ r^{-2}$ and is static ($v = 0$) at $t = 0$, while the Larson-Penston (1969) solution is $\\sim 4.4$ times denser and far from equilibrium ($v \\approx -3.3\\ \\as$). During the accretion phase ($t > 0$), the infall envelope is a factor $\\sim 7$ denser in the Larson-Penston solution. %than in the Shu model. Accordingly, the mass infall rate is also much larger in the Larson-Penston case ($\\sim 47\\, \\as^3/G$) than in the Shu case ($\\sim \\as^3/G$).  In practice, however, protostellar collapse is unlikely to be strictly self-similar, and the above similarity solutions can only be taken as plausible asymptotes. More realistic initial conditions than the SIS are provided by the so-called `Bonnor-Ebert' spheres (e.g. Bonnor~1956), which represent the equilibrium states for self-gravitating isothermal spheres and have a flat density profile in their central $\\sim $ Jeans length region. Such spheres are stable for a center-to-edge density contrast $< 14.3$ and unstable for a density contrast $> 14.3$ (e.g. Bonnor~1956). Numerical hydrodynamic simulations of cloud collapse starting from such initial conditions (e.g. Foster \\& Chevalier 1993, Hennebelle et al. 2002) find that the Larson-Penston similarity solution is generally a good approximation near point-mass formation ($t= 0$) {\\it at small radii}, but that the Shu solution is more adequate at intermediate $t \\geq 0$ times, before the expansion wave reaches the edge of the initial, pre-collapse dense core. In general, the mass accretion rate is thus expected to be time-dependent.  Observationally, it is by comparing the (density and velocity) structure of prestellar cores such as L1544 (see Fig.~\\ref{fig-intro}a) with the structure of the envelopes surrounding Class~0 protostars such as IRAM~04191 (cf. Fig.~\\ref{fig-intro}b) that one may hope to constrain the initial conditions for collapse and to discriminate between the various existing models. %In the following, we discuss these two aspects in turn. ",
        "conclusions": "In the case of low-mass, isolated dense cores, the SIS model of Shu~(1977) describes global features of the collapse (e.g. the mass infall rate) reasonably well and thus remains a useful, approximate guide. In detail, however, the extended infall velocity profiles observed in prestellar cores (e.g. Lee et al. 2001) and in the very young Class~0 object IRAM~04191 (\\S ~3.2 above) are inconsistent with the pure inside-out collapse picture of Shu (1977). The {\\it shape} of the density profiles observed in prestellar cores are well fitted by purely thermal Bonnor-Ebert sphere models, but the {\\it absolute} values of the densities are suggestive of some additional magnetic support (\\S ~2.2). The observed infall velocities are also marginally consistent with isothermal collapse models starting from %marginally stable equilibrium Bonnor-Ebert spheres (e.g. Foster \\& Chevalier 1993, Hennebelle et al. 2002), as such models tend to produce somewhat faster velocities. This suggests that the collapse of `isolated' cores is essentially {\\it spontaneous} and somehow moderated by magnetic effects in magnetized, probably not strictly isothermal Bonnor-Ebert cloudlets. Furthermore, the contrast seen in Fig.~\\ref{iram04191_model} between a steeply declining rotation velocity profile and a flat infall velocity profile beyond $\\sim 3500$~AU suggests that angular momentum is {\\it not} conserved in the outer envelope of IRAM~04191. Such a behavior is very difficult to explain in the context of non-magnetic collapse models. In the presence of a relatively strong ($\\sim 60\\, \\mu$G) magnetic field, on the other hand, the outer envelope can be coupled to, and spun down by, the (large moment of inertia of the) ambient cloud (e.g. Basu \\& Mouschovias 1994). Based on a qualitative comparison with the ambipolar diffusion models of Basu \\& Mouschovias, Belloche et al. (2002) propose that the rapidly rotating inner envelope of IRAM~04191 corresponds to a magnetically supercritical core decoupling from an environment still supported by magnetic fields and strongly affected by magnetic braking. (Quantitatively, however, the models published so far are rotating too slowly to fit the observations.) In this view, the inner $\\sim 3500$~AU radius envelope of IRAM~04191 would correspond to the effective mass reservoir ($\\sim 0.5\\, M_\\odot $) from which the central star is being built. Moreover, comparison of these results with the rotational characteristics of other %(prestellar and protostellar) objects in Taurus (Ohashi et al. 1997) suggests that IRAM~04191 behaves in a typical manner and is simply observed particularly soon after point mass formation (i.e., at $t \\simgt 0$). If this is correct, {\\it the masses of the stars forming in clouds such as Taurus may be largely determined by magnetic decoupling effects}.  \\begin{figure}[ht] %\\resizebox{\\hsize}{!}{\\includegraphics{/home/storage/pandre/tmp3/waterloo/phen_%waterloo.ps}} \\resizebox{\\hsize}{!}{\\includegraphics{andrep9_rev.ps}} \\caption{Density profiles (solid curves) obtained slightly after point mass formation ($t \\simlt 10^4$~yr) in SPH numerical simulations of the collapse of a nearly critical [$(\\rho_c/\\rho_{out})_{init} \\sim 12$] %initially stable isothermal ($T= 10$~K) Bonnor-Ebert sphere induced by external compression (Hennebelle et al. 2002). For comparison, the dotted line shows the $\\rho \\propto r^{-2}$ profile of a SIS at 10~K. In (a), the collapse was initiated quasi-statically (by very slow compression with $P_{ext}/\\dot{P}_{ext} = 20\\ \\times$ the initial sound crossing time $R_{init}/\\as $) and the density profile at $t \\simgt 0$ is very similar to that of the SIS. By contrast, in (b), the collapse was induced by a very rapid increase in external pressure (with $P_{ext}/\\dot{P}_{ext} = 0.03 \\times R_{init}/\\as$), resulting in much larger densities around $t \\sim 0$. \\label{hennebelle}} \\end{figure}  In protoclusters, by contrast, the large overdensity factors measured for Class~0 envelopes compared to hydrostatic isothermal structures (cf. \\S ~2.4), as well as the fast supersonic infall velocities and very large infall rates observed in some cases (e.g. \\S ~3.2), are inconsistent with self-initiated forms of collapse and require a {\\it strong external influence}. This point is illustrated in Fig.~\\ref{hennebelle} with the results of recent SPH simulations by Hennebelle et al. (2002). These simulations follow the evolution of a Bonnor-Ebert sphere whose collapse has been induced by an increase in external pressure $P_{ext}$. Large overdensity factors (compared to a SIS), together with supersonic infall velocities, and large infall rates ($\\simgt 10\\, \\as^3/G$) are found near $t = 0$ when (and only when) the increase in $P_{ext}$ is strong and very rapid (e.g. Fig.~\\ref{hennebelle}b), resulting in a violent compression wave. Such a violent collapse may be conducive to the formation of both massive stars (through higher accretion rates) and multiple systems (when realistic, non-isotropic compressions are considered). Future high-resolution studies with the next generation of (sub)millimeter instruments (e.g., ALMA) will greatly help test this view and shed further light on the physics of collapse in cluster-forming regions."
    },
    "0212/astro-ph0212015_arXiv.txt": {
        "abstract": "Extensive observational campaigns of afterglow hunting have greatly enriched our understanding of the gamma-ray burst (GRB) phenomenon. Efforts have been made recently to explore some afterglow properties or signatures that will be tested by the on-going or the future observational campaigns yet come. These include the properties of GRB early afterglows in the temporal domain; the GeV-TeV afterglow signatures in the spectral domain; as well as a global view about the GRB universal structured jet configuration. These recent efforts are reviewed. Within the standard cosmological fireball model, the very model(s) responsible for the GRB prompt emission is (are) not identified. These models are critically reviewed and confronted with the current data. ",
        "introduction": "Gamma-ray bursts (GRBs) belong to one of those classes of mysterious objects whose nature took great efforts to identify. For years after their discovery, our knowledge about the objects had been limited to the ``gamma-ray'' in the spectral domain, and the ``burst'' in the temporal domain. The revolutionary discovery and the extensive hunting and monitoring of the broadband afterglows for dozens of GRBs greatly extended our knowledge about the events both in the spectral domain (from radio to X-rays) and in the temporal domain (from minutes to years). We now have some good knowledge about at least one subset of GRBs, i.e. the so-called long GRBs whose durations are longer than 2 seconds. These include that they are originated from cosmological distances, likely associated with the deaths of massive stars; that the afterglow is emission from the external forward shock when the fireball impacts the ambient interstellar medium (ISM); and that the fireballs are likely collimated at least for some GRBs (e.g., Piran 1999; van Paradjis, Kouveliotou, \\& Wijers 2000; M\\'esz\\'aros 2002). It is expected that some new progress will be made in the GRB study in the coming years, thanks to the observational advances led by some future space and ground-based multi-wavelength telescopes, especially by the two NASA missions, Swift and GLAST. Among others, the very early afterglows including the bridge between the GRB prompt phase and the afterglow phase will be carefully studied; the broad-band afterglow campaign will be extended to high energy (GeV-TeV) regimes; and a large sample of dataset combining the redshifts and the GRB prompt emission \\& afterglow properties will be available, leading to a global view of the GRB jet configuration and the identification of the GRB prompt emission site(s) and mechanism(s). Below we will review some recent theoretical efforts in anticipation of these upcoming observational breakthroughs. ",
        "conclusions": "A new epoch for the GRB study is coming within the next several years, especially following the launches of Swift and GLAST, and the utilization of some other broad-band advanced facilities. It is optimistic to expect the following breakthroughs in the coming years:  \\ref $\\bullet$ Early afterglows will be carefully studied. The missing link between the prompt emission and the afterglow will be identified, including the detailed reverse shock information. The information from the central engine may be retrievable through well-modeled injection signatures. \\ref $\\bullet$ High energy afterglows will be monitored and studied in conjunction with the low energy afterglows, which will bring invaluable information about the unknown shock physics and the GRB environments. \\ref $\\bullet$ The GRB jet configuration will be identified. We expect the universal structured jet model will be validated by future data. \\ref $\\bullet$ With accumulation of a large sample of the redshift and spectral information for GRBs/XRFs in the Swift era, the right emission site(s) and mechanism(s) for the prompt emission may be identified (or at least strongly constrained). The fireball content and the nature of the relativistic wind from the central engine may be also understood."
    },
    "0212/astro-ph0212223_arXiv.txt": {
        "abstract": "This article describes the computation of cosmic microwave background anisotropies in a universe with multi-connected spatial sections and focuses on the implementation of the topology in standard CMB computer codes.  The key ingredient is the computation of the eigenmodes of the Laplacian with boundary conditions compatible with multi-connected space topology. The correlators of the coefficients of the decomposition of the temperature fluctuation in spherical harmonics are computed and examples are given for spatially flat spaces and one family of spherical spaces, namely the lens spaces.  Under the hypothesis of Gaussian initial conditions, these correlators encode all the topological information of the CMB and suffice to simulate CMB maps.  \\vspace*{0.35cm} \\noindent {\\footnotesize Preprint numbers: SPhT-T02/182, LPT-02/123, {\\tt astro-ph/0212223}} ",
        "introduction": "\\label{sec_intro}  Future cosmic microwave background (CMB) experiments such as the MAP~\\cite{map} and later the Planck satellites~\\cite{planck} will provide full sky maps of CMB anisotropies (up to the galactic cut). These datasets offer the opportunity to probe the topological properties of our universe. A series of tests to detect the topology, including the use of the angular power spectrum~\\cite{sokolov93,staro93,stevens93,costa95}, the distribution of matched patterns such as circles~\\cite{cornish98}, correlation of antipodal points~\\cite{levin98} and non Gaussianity~\\cite{inoue00} have been proposed (see Refs.~\\cite{lachieze95,uzan99,levin02} for reviews of CMB methods to search for the topology). The study of the detectability of the topology by any of these methods first requires simulating maps with the topological signature for a large set of topologies.  These maps will allow one to test the detection methods, estimate their run time, and, once all sources of noise are added, determine to what extent a given method detects the topological signal (in the same spirit as the investigation of the ``crystallographic'' methods based on galaxy catalogs~\\cite{us3}).  In a simply connected\\footnote{Geometers and cosmologists often refer to simple-connectedness as the ``trivial topology''. However, trivial topology has a different meaning in the context of point-set topology: in that formalism, the trivial topology is the smallest topology on a set $X$, namely the one in which the only open sets are the empty set and the entire set $X$.}, spatially homogeneous and isotropic universe, the angular correlation function depends only on the angle between the two directions and the coefficients $\\ALM{\\ell}{m}$ of the decomposition of the temperature fluctuation in spherical harmonics, which are uncorrelated for different sets of $\\ell$ and $m$. Multi-connectedness breaks the global isotropy and sometimes the global homogeneity of the universe, except in projective space (see, e.g., Ref.~\\cite{lwugl}). Consequently, the CMB temperature angular correlation function will depend on the two directions of observation, not only on their relative angle, and possibly on the position of the observer. This induces correlations between the $\\ALM{\\ell}{m}$ of different $\\ell$ and $m$. Such correlations are hidden when one considers only the angular correlation function and its coefficients, the so-called $C_\\ell$, in a Legendre polynomial decomposition, because they pick up only the isotropic part of the information and are therefore a poor indicator of the topology. This work aims to detail the whole computation of the correlation matrix \\begin{equation} \\label{eq:I:1} \\CLMLPMP{\\ell}{m}{\\ell'}{m'} \\equiv \\left< \\ALM{\\ell}{m} \\ALM{\\ell'}{m'}^* \\right> , \\end{equation} which encodes all the topological properties of the CMB, and from which one can compute the usual $C_\\ell$, simulate maps, and so on.  The study of the detectability of a topological signal (if it exists) in forthcoming CMB datasets requires simulating high quality maps containing the topological signature for a wide class of topologies. Up to now, most CMB studies considered only compact Euclidean spaces~\\cite{sokolov93,staro93,stevens93,costa95,levin99,inoue01} and some compact hyperbolic spaces~\\cite{aurich00,inoue00b,cornish00,bond1,bond2,bond3}, and focused mainly on the $C_\\ell$. The approach developed in this article, and first introduced in Ref.~\\cite{lwugl}, is well suited to simulate the required CMB maps in any topology once the eigenmodes of the Laplacian have been determined. It paves the way to the simulation of maps for a wide range of topologies, particularly spherical ones.  Recent measurements of the density parameter $\\Omega$ imply that the observable universe is ``approximately flat''\\footnote{The popular expression ``flat universe'' is misleading, because in General Relativity the ``universe'' is not three-, but four-dimensional and the Friedmann-Lema\\^{\\i}tre solutions are dynamical, so that the universe is not flat --- only its spatial section may be flat or nearly so. In what follows, we will implicitly assume we are talking about the {\\em three-dimensional spacelike sections} of the universe when talking about flat, hyperbolic, or spherical spaces/universes.}, perhaps with a slight curvature. The exact constraint on the total density parameter obtained from CMB experiments depends on the priors used during the data analysis. For example with a prior on the nature of the initial conditions, the Hubble parameter and the age of the universe, recent analysis of the DASI, BOOMERanG, MAXIMA and DMR data~\\cite{sievers,netterfield,archeops2} lead to $\\Omega = 0.99 \\pm 0.12$ at $1\\sigma$ level, and to $\\Omega = 1.04 \\pm 0.05$ at $1\\sigma$ level if one takes into account only the DASI, BOOMERanG and CBI data.  Including stronger priors can indeed sharpen the bound. For instance, including information respectively on large scale structure and on supernovae data leads to $\\Omega = 1.01_{- 0.06}^{+ 0.09}$ and $\\Omega = 1.02_{- 0.08}^{+ 0.09}$ at $1\\sigma$ level while including both finally leads to $\\Omega = 1.00_{- 0.06}^{+ 0.10}$. This has been recently improved by the Archeops balloon experiments~\\cite{archeops2,archeops1} which get, with a prior on the Hubble constant, $\\Omega = 1.00_{- 0.02}^{+ 0.03}$. In conclusion, it is fair to assert that current cosmological observations set a reliable bound $0.9 < \\Omega < 1.1$. These results are consistent with Friedmann-Lema\\^{\\i}tre universe models with spherical, flat or hyperbolic spatial sections.  In the spherical and hyperbolic cases, $\\Omega \\simeq 1$ implies that the curvature radius must be larger than the horizon radius. In all three cases --- spherical, flat, and hyperbolic --- the universe may be simply connected or multiply connected.\\\\  The possibility of detecting the topology of a nearly flat universe was discussed in Ref.~\\cite{wlu02}. It was noted that the chances of detecting a multiply connected topology are worst in a large hyperbolic universe. The reason is that the typical translation distance between a cosmic source and its nearest topological image seems to be on the order of the curvature radius, and that when $\\Omega \\simeq 1$ the distance to the last scattering surface is less than the half of that distance.  See Refs.~\\cite{gomero1,aurichsteiner,inoue3,gomero3,weeks} for some studies on detectability of nearly flat hyperbolic universes. In a multiply connected flat universe the topology scale is completely independent of the horizon radius, because Euclidean geometry --- unlike spherical and hyperbolic geometry --- has no preferred scale and admits similarities.  Note that in the Euclidean case, there are only ten compact topologies, which reduces and simplifies the analysis (in particular regarding the computation of the eigenmodes of the Laplacian). In a spherical universe the topology scale depends on the curvature radius, but, in contrast to the hyperbolic case, as the topology of a spherical manifold gets more complicated, the typical distance between two images of a single cosmic source decreases. No matter how close $\\Omega$ is to $1$, only a finite number of spherical topologies are excluded from detection. The particular case of the detectability of lens spaces was studied in Ref.~\\cite{gomero2}, which also considers the detectability of hyperbolic topologies.\\\\  At present, the status of the constraints on the topology of the universe is still very preliminary.  Regarding locally Euclidean spaces, it was shown on the basis of the COBE data that in the case of a vanishing cosmological constant the size of the fundamental domain of a 3-torus has to be larger than $L \\geq 4800 \\,h ^{-1} \\UUNIT{Mpc}{}$~\\cite{sokolov93,staro93,stevens93,costa95}, where the length $L$ is related to the smallest wavenumber $2 \\pi / L$ of the fundamental domain, which induces a suppression of fluctuations on scales beyond the size $L$ of the fundamental domain.  This constraint does not exclude a toroidal universe since there can be up to eight copies of the fundamental cell within our horizon.  This constraint relies mainly on the fact that the smallest wavenumber is $2 \\pi / L$, which induces a suppression of fluctuations on scales beyond the size of the fundamental domain.  This result was generalized to all Euclidean manifolds in Ref.~\\cite{levin99}.  A non-vanishing cosmological constant induces more power on large scales, via the integrated Sachs-Wolfe effect.  For instance if $\\Omega_\\Lambda = 0.9$ and $\\Omega_\\MAT = 0.1$, the constraint is relaxed to allow for 49 copies of the fundamental cell within our horizon.  The constraint is also milder in the case of compact hyperbolic manifolds and it was shown~\\cite{aurich00,inoue00b,cornish00} that the angular power spectrum was consistent with the COBE data on multipoles ranging from 2 to 20 for the Weeks and Thurston manifolds.  Another approach, based on the method of images, was developed in~\\cite{bond1,bond2,bond3}. Only one spherical space using this method of images was considered in the literature, namely projective space~\\cite{souradeep}. The tools developed in this article, as well as in our preceding works~\\cite{lwugl,glluw,luw02}, will let us fill the gap regarding the simulation of CMB maps in spherical universes.  Technically, in standard relativistic cosmology, the universe is described by a Friedmann-Lema\\^{\\i}tre spacetime with locally isotropic and homogeneous spatial sections.  These spatial sections can be defined as constant density or time hypersurfaces.  With such a splitting, the equations of evolution of the cosmological perturbations, which give birth to the large scale structures of the universe, reduce to a set of coupled differential equations involving the Laplacian.  This system is conveniently solved in Fourier space. In the case of a multiply connected universe, we visualize space as the quotient $X / \\Gamma$ of a simply connected space $X$ (which is just a $3$-sphere $\\STR$, a Euclidean space $\\RTR$, or a hyperbolic space $\\HTR$, depending on the curvature) by a group $\\Gamma$ of symmetries of $X$ that is discrete and fixed point free.  The group $\\Gamma$ is called the holonomy group.  To solve the evolution equations we must first determine the eigenmodes $\\UPSTKS{\\Gamma}{k}{}$ and eigenvalues $- \\kappa_k^2$ of the Laplacian on $X / \\Gamma$ through the generalized Helmholtz equation \\begin{equation} \\label{Helmotz1} \\Delta \\UPSTKS{\\Gamma}{k}{} = -\\kappa_k^2 \\UPSTKS{\\Gamma}{k}{} , \\end{equation} with \\begin{equation} \\kappa_k^2 = k^2 - K , \\end{equation} where $k$ indexes the set of eigenmodes, the constant $K$ is positive, zero or negative according to whether the space is spherical, flat or hyperbolic, respectively,\\footnote{We work here in comoving units, but when spatial section are not flat, the curvature of the comoving space is {\\em not} normalized: it is not assumed to be $+ 1$ in the spherical case nor is it assumed to be $- 1$ in the hyperbolic case. Thus our eigenvalues may differ numerically from those found in the mathematical literature, although of course they agree up to a fixed constant multiple $|K|^{-\\frac{1}{2}}$.} and the boundary conditions are compatible with the given topology. The Laplacian in Eq.~(\\ref{Helmotz1}) is defined as $\\Delta \\equiv D^i D_i$, $D_i$ being the covariant derivative associated with the metric $\\gamma_{i j}$ of the spatial sections ($i$, $j = 1$, $2$, $3$).  The eigenmodes of $X / \\Gamma$, on which any function on $X / \\Gamma$ can be developed, respect the boundary conditions imposed by the topology. That is, the eigenmodes of $X / \\Gamma$ correspond precisely to those eigenmodes of $X$ that are invariant under the action of the holonomy group $\\Gamma$.  Thus any linear combination of such eigenmodes will satisfy, by construction, the required boundary conditions.  In this way we visualize the space of eigenmodes of $X / \\Gamma$ as a subspace of the space of eigenmodes of $X$, namely the subspace that is invariant under the action of $\\Gamma$.  The computational challenge is to find this $\\Gamma$-invariant subspace and construct an orthonormal basis for it.  In the case of flat manifolds the eigenmodes of $X / \\Gamma = \\RTR / \\Gamma$ can be found analytically.  In the case of hyperbolic manifolds, many numerical investigations of the eigenmodes of $X / \\Gamma = \\HTR / \\Gamma$ have been performed~\\cite{inoue00b,aurich89a,aurich89b,inoue99,aurich96,cornish99}. In the case of spherical manifolds, the eigenmodes of $X / \\Gamma = \\STR / \\Gamma$ have been found analytically for lens and prism spaces~\\cite{luw02} and otherwise can be found numerically~\\cite{lwugl}.  In the following, we will develop the eigenmodes of $X / \\Gamma$ on the basis $\\YXKLM{X}{k}{\\ell}{m}$ of the eigenmodes of the universal covering space as \\begin{equation} \\label{eq:1} \\UPSTKS{\\Gamma}{k}{s} = \\sum_{\\ell = 0}^\\infty \\sum_{m = -\\ell}^{\\ell} \\XITKSLM{\\Gamma}{k}{s}{\\ell}{m} \\YXKLM{X}{k}{\\ell}{m} , \\end{equation} so that all the topological information is encoded in the coefficients $\\XITKSLM{\\Gamma}{k}{s}{\\ell}{m}$, where $s$ labels the various eigenmodes sharing the same eigenvalue $-\\kappa_k^2$, both of which are discrete numbers\\footnote{The spectrum is discrete when the space is compact, e.g. the torus or any spherical space. In a non-compact multi-connected space such as a cylinder it will have a continuous component.}.  The sum over $\\ell$ runs from $0$ to infinity if the universal covering space is non-compact (i.e., hyperbolic or Euclidean). The $\\XITKSLM{\\Gamma}{k}{s}{\\ell}{m}$ coefficients can be determined analytically for Euclidean manifolds. In the case of a spherical manifold, the modes are discrete, \\begin{equation} \\label{knu} k = (\\nu + 1) \\sqrt{K} \\end{equation} where $\\nu$ is a nonnegative integer (the cases $\\nu = 0$ and $\\nu = 1$ correspond to pure gauge modes~\\cite{abott86}), so that $\\kappa_k^2 = \\nu (\\nu + 2) K$, and the sum over $\\ell$ is finite and runs from 0 to $\\nu$ since the multiplicity of a mode $k$ is at most its multiplicity in the universal cover, which is $(\\nu + 1)^2$. Among spherical spaces, the case of prism and lens spaces are the simplest since one can determine these coefficients analytically~\\cite{luw02}. The worst situation is that of compact hyperbolic manifolds, which is analogous to the Euclidean case since the universal covering $\\RTR$ is not compact but for which no analytical forms are known for the eigenmodes. One then needs to rely on numerical computations (see, e.g., Refs.~\\cite{cornish99,glluw}).  Our preceding works provided a full classification of spherical manifolds~\\cite{glluw} and developed methods to compute the eigenmodes of the Laplacian in them~\\cite{lwugl}. Among all spherical manifolds, we were able to obtain analytically the spectrum of the Laplacian for lens and prism spaces~\\cite{luw02} which form two countable families of spherical manifolds. The goal of the present article is to simulate CMB maps for these two families of spherical topologies, as well as for the Euclidean topologies.  We first review briefly the physics of CMB anisotropies (Sec.~\\ref{sec_cmb}) mainly to explain how to take into account the topology (Sec.~\\ref{subsec_implement}) once the coefficients $\\XITKSLM{\\Gamma}{k}{s}{\\ell}{m}$ are determined. We then detail the computation of these coefficients, focusing on the cases where it can be performed analytically, that is for flat spaces and lens and prism spaces.  We then present results of numerical simulations (Sec.~\\ref{sec_simulation}) as well as simulated maps. We discuss these maps qualitatively and confirm that the expected topological correlations (namely matching circles~\\cite{cornish98}) are indeed present. The effects of the integrated Sachs-Wolfe and Doppler terms, as well as the thickness of the last scattering surface, are discussed in order to give some insight into the possible detectability of these correlations. Figure~\\ref{method} summarizes the different independent steps of the computation as well as their interplay.  \\section*{Notations}  The local geometry of the universe is described by a Friedmann-Lema\\^{\\i}tre metric \\begin{equation} \\label{fl_metric} \\ddd s^2 = - c^2 \\ddd t^2 + a^2 (t) \\left[\\ddd \\CHU^2 + \\SK{K}^2 (\\CHU) \\ddd \\Omega^2 \\right] , \\end{equation} where $a(t)$ is the scale factor, $t$ the cosmic time, $\\ddd \\Omega^2 \\equiv \\ddd \\theta^2 + \\sin^2 \\theta \\, \\ddd \\varphi^2$ the infinitesimal solid angle, $\\CHU$ is the comoving radial distance, and where \\begin{equation} \\label{def_sk} \\SK{K}(\\CHU) = \\left\\lbrace \\begin{array}{ll} \\sinh (\\sqrt{|K|}\\;\\CHU) / \\sqrt{|K|} & \\quad {\\rm (hyperbolic \\, case)} \\\\ \\CHU                                  & \\quad {\\rm (flat \\, case)      } \\\\ \\sin (\\sqrt{K}\\;\\CHU) / \\sqrt{K}      & \\quad {\\rm (spherical \\, case) } \\end{array} \\right. \\end{equation} In the case of spherical and hyperbolic spatial sections, we also introduce the dimensionless coordinate \\begin{equation}\\label{def_chn} \\CHN=\\sqrt{|K|}\\CHU, \\end{equation} which expresses radial distances in units of the curvature radius.  \\begin{center} \\begin{figure} \\unitlength=1cm \\begin{picture}(18,18) \\thicklines \\put(0,0) {\\framebox(9.5,8){}} \\put(0.25,7.5) {\\SZFIG III: Eigenmodes of $X / \\Gamma$} \\put(0.5,6.75) {\\circle{0.8} \\makebox(0,0) {\\SZFIG$\\Gamma$}} \\put(0.9,6.7) {\\vector(1,0){3}} \\thinlines \\put(4,6.35) {\\framebox(4.25,.8){}} \\thicklines \\put(4.25,6.6) {\\SZFIG $\\UPSTKS{\\Gamma}{k}{s} = \\sum_{\\ell, m} \\XITKSLM{\\Gamma}{k}{s}{\\ell}{m} \\YXKLM{X}{k}{\\ell}{m}$} \\put(0.5,3.25) {\\circle{0.8} \\makebox(0,0) {\\SZFIG$K$}} \\put(0.7,3.6) {\\vector(2,3){0.9}} \\put(0.9,3.2) {\\vector(1,0){0.7}} \\put(0.7,2.8) {\\vector(2,-3){0.9}} \\put(1.6,5) {\\dashbox{0.05}(0.8,0.8) {\\SZFIG$< 0$}} \\put(2.75,5.25) {\\SZFIG numerical~\\cite{cornish99,glluw}} \\put(1.6,2.85) {\\dashbox{0.05}(0.8,0.8) {\\SZFIG$= 0$}} \\put(2.75,3.1) {\\SZFIG analytical~\\cite{lachieze95,levin02}\\ldots} \\put(1.6,0.6) {\\dashbox{0.05}(0.8,0.8) {\\SZFIG$> 0$}} \\put(2.75,1.2) {\\SZFIG general case: numerical~\\cite{glluw}} \\put(2.75,0.5) {\\SZFIG lens and prism: analytical~\\cite{luw02,glluw}} \\put(7.1,0.15) {(see Sec.~\\ref{subsubsec_lens})} \\put(0,10) {\\framebox(9.5,8){}} \\put(0.25,17.5) {\\SZFIG I: Universal cover ($X$) eigenmodes} \\thinlines \\put(0.1,16.5) {\\framebox(9.25,.75){}} \\thicklines \\put(0.25,16.75) {$\\Delta \\YXKLM{X}{k}{\\ell}{m} = - (k^2 - K) \\YXKLM{X}{k}{\\ell}{m} \\;;\\quad \\YXKLM{X}{k}{\\ell}{m} = \\RXKL{X}{k}{\\ell} (\\CHU) \\YLM{\\ell}{m} (\\theta, \\varphi)$} \\put(0.5,13.25) {\\circle{0.8} \\makebox(0,0){$K$}} \\put(0.7,13.6) {\\vector(2,3){0.9}} \\put(0.9,13.2) {\\vector(1,0){0.7}} \\put(0.7,12.8) {\\vector(2,-3){0.9}} \\put(1.6,15) {\\dashbox{0.05}(0.8,0.8) {$< 0$}} \\put(2.75,15.25) {$X = \\HTR:\\, \\RXKL{X}{k}{\\ell} = \\sqrt{\\frac{N_{k \\ell}}{k \\SK{K} (\\CHU)}} \\PLM{- 1 / 2 + i \\omega}{-1 / 2 - \\ell} \\left(\\cosh\\CHN \\right)$} \\put(4.25,14.6) {$\\omega \\in[0,\\infty[\\quad\\hbox{or}\\quad i \\omega \\in[0,1]$} \\put(1.6,12.85) {\\dashbox{0.05}(0.8,0.8) {$0$}} \\put(2.75,13.1) {$X = \\RTR:\\, \\RXKL{X}{k}{\\ell} = \\sqrt{\\frac{2}{\\pi}} j_\\ell(k \\CHU)$} \\put(4.25,12.45) {$k\\in[0,\\infty[$} \\put(1.6,10.6) {\\dashbox{0.05}(0.8,0.8) {$> 0$}} \\put(2.75,10.85) {$X = \\STR:\\, \\RXKL{X}{k}{\\ell} = \\sqrt{\\frac{M_{k \\ell}}{k \\SK{K} (\\CHU)}} \\PLM{1 / 2 + \\nu}{- 1 / 2 - \\ell} \\left(\\cos\\CHN \\right)$} \\put(4.25,10.2) {$\\nu=2,3,\\ldots$} \\put(7.5,10.15) {(see  App.~\\ref{app_A})} \\put(11,8.5) {\\framebox(7,1){}} \\put(11.25,8.9) {\\SZFIG V: Output: $\\ALM{\\ell}{m}$, maps, $C_\\ell$\\ldots} \\put(16.1,8.65) {(see Sec.~\\ref{sec_simulation})} \\put(11,10) {\\framebox(7,8){}} \\put(11.25,17.5) {\\SZFIG II: CMB in $X$} \\thinlines \\put(11.25,16.5) {Perturbation equations} \\put(11.8,16.2) {(local physics)} \\put(15.25,16.5) {Initial conditions} \\put(12,15) {\\SZFIG $O_k^{[X]} \\left(\\RXKL{X}{k}{\\ell}\\right)$} \\put(12.25,14) {\\SZFIG $G_\\ell(k)$} \\put(15.9,14) {\\SZFIG $\\PS (k)$} \\put(12.5,12) {\\SZFIG $\\left<\\ALM{\\ell}{m} \\ALM{\\ell'}{m'}^* \\right> = C_\\ell \\KRON{\\ell}{\\ell'} \\KRON{m}{m'}, $} \\put(12.5,11) {\\SZFIG $C_\\ell = \\frac{2}{\\pi} \\int k^2 \\ddd k G_\\ell^2(k) \\PS (k)$} \\put(15.9,10.15) {(see Sec.~\\ref{subsec_local})} \\thicklines \\put(12.25,10.75) {\\dashbox{0.05}(4.75,1.75){}} \\put(12.8,16){\\vector(0,-1){.6}} \\put(12.8,14.8){\\vector(0,-1){.5}} \\put(16.2,16.3){\\vector(0,-1){2}} \\put(13.2,14.15){\\framebox(2.6,0){}} \\put(14.5,14.15){\\vector(0,-1){1.6}}  \\put(11,0) {\\framebox(7,8){}} \\put(11.25,7.5) {\\SZFIG IV: CMB in $X / \\Gamma$} \\put(15.9,0.15) {(see Sec.~\\ref{subsec_implement})} \\put(12.75,6.5) {\\SZFIG $\\left<\\ALM{\\ell}{m} \\ALM{\\ell'}{m'}^* \\right> = \\CLMLPMP{\\ell}{m}{\\ell'}{m'} ,$} \\put(11.75,5.5) {\\SZFIG $\\CLMLPMP{\\ell}{m}{\\ell'}{m'} = \\frac{2}{\\pi} \\sum_{k, s} \\XITKSLM{\\Gamma}{k}{s}{\\ell}{m} \\XITKSLM{\\Gamma}{k}{s \\; *}{\\ell'}{m'}$} \\put(14.14,4.9) {\\SZFIG $\\quad G_\\ell(k) G_{\\ell'}(k) \\PS (k)$} \\put(11.6,4.65) {\\dashbox{0.05}(5.9,2.4){}} \\thinlines \\put(11.4,1) {Perturbation equations} \\put(11.15,0.7) {(local physics, unchanged)} \\put(15.25,1) {Initial conditions} \\put(15.5,0.7) {(unchanged)} \\put(12,2) {\\SZFIG $O_k^{[X]} \\left(\\UPSTKS{\\Gamma}{k}{s}\\right)$} \\put(12.25,3.5) {\\SZFIG $G_\\ell(k)$} \\put(15.9,3.5) {\\SZFIG $\\PS (k)$} \\thicklines \\put(12.7,1.3){\\vector(0,1){.6}} \\put(12.7,2.4){\\vector(0,1){1}} \\put(16.2,1.3){\\vector(0,1){2.05}} \\put(13.2,3.6){\\framebox(2.6,0){}} \\put(14.5,3.6){\\vector(0,1){1.0}} \\put(4.75,10) {\\vector(0,-1){2}} \\put(9.5,14) {\\vector(1,0){1.5}} \\put(9.5,4) {\\vector(1,0){1.5}} \\put(14.5,10) {\\vector(0,-1){0.5}} \\put(14.5,8) {\\vector(0,1){0.5}} \\end{picture} \\caption{The computation of CMB anisotropies in multi-connected spaces follows a series of independent steps. From the knowledge of the spatial curvature $K$, one determines the universal covering space, $X$, as well as the eigenmodes of the Laplacian in this simply-connected space (upper left box); these functions are well-known and frequently used in standard cosmology computations, and are recalled in Appendix~\\ref{app_A} [we have introduced $\\omega = k / \\sqrt{|K|}$ for the hyperbolic case and $\\nu = k / \\sqrt{K} - 1$ for the spherical case [recall Eq.~(\\ref{knu})], and used $\\CHN$ defined by Eq.~(\\ref{def_chn})]. Once some cosmological parameters and a scenario of structure formation has been chosen (upper right box), the CMB anisotropies in the universal covering space can be computed. This step is also a standard step that can be achieved numerically by a number of codes (see Sec.~\\ref{subsec_local}). An independent computation (lower left box) is the determination of the eigenmodes of the Laplacian compatible with the topology of space, specified by the choice of the holonomy group $\\Gamma$. Our approach is to encode all the topological information in a set of parameters $\\XITKSLM{\\Gamma}{k}{s}{\\ell}{m}$. Their computation is described in Sec.~\\ref{subsubsec_lens} and can be performed either numerically or analytically according to the case at hand. The implementation of the topology in a standard CMB code (lower right box) is described in Sec.~\\ref{subsec_implement} and yields the complete correlation matrix of the $\\ALM{\\ell}{m}$ from which one can (center right box) compute $C_\\ell$, simulate maps, etc.} \\label{method} \\end{figure} \\end{center} ",
        "conclusions": "\\label{sec_conclusion}  This articles describes the implementation of topology in CMB codes and gives explicitly the required tools to perform such an implementation in flat and spherical spaces.  As emphasized in the Introduction, these two cases are observationally the most relevant for an almost flat universe.  Examples of simulated maps were given in the two cases. Here we presented only low resolution maps due to the computational time limitation but higher resolution maps will be presented elsewhere.  It was checked that the expected topological correlations (the matched circles) were present, confirming the quality of our simulations.  Our method relies on the computation of the correlation matrix of the coefficients of the decomposition of the temperature fluctuation in spherical harmonics. This matrix encodes all the topological information. We emphasize that, due to the breakdown of global isotropy, this matrix is not purely diagonal. This also offers a working example to construct tests for the detection of deviation from global isotropy.  We have illustrated the influence of different effects that will tend to blur these patterns and affect the perfect circle matching, namely the Doppler effect and the integrated Sachs-Wolfe effect.  We also considered the effect of the thickness of the last scattering surface, but found it to be negligible on the scales considered here. A more detailed quantitative analysis of these effects on the detectability of the topological signal is left for future studies~\\cite{future}.  A complete investigation of the detectability of the topology in coming CMB data requires the construction of reliable simulation tools. Besides the quantification of the amplitude of the effects cited above, one would also need to include all other observational effects such as instrumental noise, foreground contamination, etc. The present work paves the way to all these essential studies."
    },
    "0212/astro-ph0212545_arXiv.txt": {
        "abstract": "Recent {\\it Hubble Space Telescope} photometry in the nearby elliptical galaxy NGC 5128 shows that its halo field star population is dominated by moderately metal-rich stars, with a peak at [m/H] $\\simeq$ $-0.4$ and with a very small fraction of metal-poor ([m/H] $<$ $-1.0$) stars.  In order to investigate the physical processes which may have produced this metallicity distribution function (MDF), we consider a model in which NGC 5128 is formed by merging of two major spiral galaxies. We find that the halo of an elliptical formed this way is predominantly populated by moderately metal-rich stars with [m/H] $\\sim$ $-0.4$ which were initially within the outer parts of the two merging discs and were tidally stripped during the merger. To match the NGC 5128 data, we find that the progenitor spiral discs must have rather steep metallicity gradients similar to the one defined by the Milky Way open clusters, as well as sparse metal-poor haloes (5\\% or less of the disc mass).  Very few stars from the central bulges of the spiral galaxies end up in the halo, so the results are not sensitive to the relative sizes (bulge-to-disc ratios) or metallicities of the initial bulges. Finally, we discuss the effects on the globular cluster system (GCS). The emergent elliptical will end up with metal-poor halo clusters from the original spiral haloes, but with moderately metal-rich halo stars from the progenitor discs, thus creating a mean offset between the MDFs of the halo stars and the GCS. Remaining questions yet to be answered concern the total size of the GCS population (the ``specific frequency problem'') and the observed existence of metal-rich globular clusters in large numbers in the outer haloes of giant ellipticals. We also discuss possible differences in the MDFs of stellar haloes of galaxies of different Hubble type. ",
        "introduction": "Physical properties of metal-deficient stellar haloes in galaxies are considered to provide vital clues to the understanding of early dynamical and  chemical evolution   of galaxies. In particular, the detailed investigation of structural, kinematic, and chemical properties of such  a ``fossil record''  component (i.e., stellar halo) of the Galaxy has revealed a possible scenario as to  how the Galaxy has developed its dynamical structures such as bulge, thin and thick discs (e.g., Freeman 1987; Majewski 1993; van den Bergh 1996). Because of  the observational difficulties in revealing three dimensional structure and kinematics  of halo stars  (with respect to their host galaxy) in each of  the Local Group galaxies other than the Milky Way, the metallicity distribution function (hereafter MDF)  of halo stars in these galaxies has served to give some constraints on the past star formation and chemical evolution histories of these systems (Mould \\& Kristian 1986; Durrell, Harris, \\& Pritchet 1994, 2001; Pritchet \\& van den Bergh 1998; Christian \\& Heasley 1986: Reitzel, Guhathakurta, \\& Gould 1998; Grillmair et al. 1996; Han et al. 1997; Martinez-Delagado \\& Aparicio 1998). The stellar halo of M31 appears to be dominated by a moderately high-metallicity population with [m/H]  $\\sim$ $-0.5$ (cf.~the references cited above), a strikingly different enrichment level than in the Milky Way stellar halo, and a strong indicator of a different formation history (e.g.~Durrell, Harris, \\& Pritchet 2001).  Although there have been many MDF studies of halo stars in Local Group members,  galaxies beyond the Local Group have only recently been investigated through {\\it HST}/WFPC2 photometry and for only a small number of cases:  the giant E/S0 NGC 5128 (Soria et al. 1996; Harris, Harris, \\& Poole 1999 hereafter referred to as HHP; Harris \\& Harris 2000, HH00; Harris \\& Harris 2002, HH02; Marleau et al. 2000), the edge-on S0 NGC 3115 (Elson 1997; Kundu \\& Whitmore 1998), and two dwarf ellipticals in the M81 group (Caldwell et al. 1998). Furthermore, Rejkuba et al. (2002) have demonstrated the existence of a significant intermediate-age AGB population in parts of the NGC 5128 halo, based on deep $U$, $V$, and $K_{s}$ color-magnitude and color-color diagrams of halo stars resolved by VLT.  A remarkable result emerging from the {\\it HST}/WFPC2 photometry of the old-halo red giant stars in NGC 5128 is the dominance of moderately metal-rich stars in the range $-1$ $<$ [m/H] $<$ 0.0 (HHP, HH00, and HH02).  HH00 and HH02  suggested that the relatively high mean abundance and small fraction ($\\sim$ 10\\%) of metal-poor stars with [m/H] $<$ $-1$ can give strong constraints on the total mass of dwarf galaxies that could be accreted by NGC 5128 and then disrupted to form its faint stellar halo.  To date, however, there are only a few attempts at physical interpretation of this material.  HH00 discussed a one-zone model of chemical evolution of this galaxy and suggested that it experienced two fairly distinct stages of halo formation: an early ``accreting box'' stage when the galaxy is assembling  through  infall  of star-forming gas clumps, and then a major stage of ``closed-box'' evolution when the star formation proceeds in the galaxy with little gas infall. HH02 extended this basic picture further to a more generalized accreting-box formation model, in which the infall rate of unenriched gas is envisaged to start at a high rate and then dies away exponentially, while star formation continues throughout until the gas supply is exhausted.  With appropriate choices of gas infall rate and the effective chemical yield, excellent overall matches to the observed MDFs can be obtained. Recently, Beasley et al. (2002a, b) have further discussed the origin of the stellar halo MDF within the context of a semi-analytic model of galaxy formation (GALFORM), based on the Cold Dark Matter (CDM) picture.  In this model, star formation can take place in two modes, a ``quiescent'' mode within the relatively unenriched pregalactic clouds, and a ``starburst'' mode whenever two roughly equal clouds suddenly merge.  They demonstrated that the global metallicity distribution within model galaxies comparable to NGC 5128 in size is broadly consistent with the observed MDF, although the model is unable to provide any information on the spatial distribution or radial change in the MDF.  These previous studies are based on one-zone models of chemical evolution and thus cannot spatially resolve the various distinct stellar components (i.e., halo, bulge,  and disc) in  a galaxy. Accordingly, comparisons of these models with the MDFs observed at different places in NGC 5128 have limited value. Numerical simulations which enable us to resolve spatially the halo, bulge, and disc components and thereby to investigate the MDF for each of these components should be very helpful for better understanding the formation of galactic stellar haloes.   The purpose of this paper is to investigate the origin of the MDF of the stellar halo of NGC 5128 based on numerical simulations of elliptical galaxy formation. We here adopt the basic ``merger'' approach (Toomre 1977) in which elliptical galaxies are proposed to be formed by major merging of two spiral galaxies.  Within the context of this assumption, we investigate the final MDF of the outer halo component of the merger remnant (i.e., the giant E galaxy) as well as its dependence on the input parameters of the progenitor discs (e.g., the MDFs of discs, bulge-to-disc-ratio, and halo mass fraction of the disc). Comparing the simulated MDFs with the observed ones for NGC 5128, we discuss the following points: (1) how the dynamics of galaxy merging is important for the formation of the gE stellar halo, (2) how the initial MDF of a disc (or bulge) in a merger controls the final MDF of the stellar halo, (3) whether the bulge-to-disc ratio of a merger progenitor spiral is important for the determination of the stellar halo's MDF, (4) what initial conditions of galaxy mergers in the present simulations can best give the MDF similar to the observed one, (5) the relevance of the MDF for the globular cluster system, which is different from the field-halo stars, and (6) whether, on the grounds of the merger picture, we should always expect the MDFs of stellar haloes in elliptical galaxies (E) to be systematically different (more metal-rich) from those in  late-type spirals (Sp).  The plan of the paper is as follows: In the next section, we describe our  numerical models  for stellar halo formation in galaxy mergers. In \\S 3, we present the numerical results mainly on the final MDFs of  merger remnants (i.e., elliptical galaxies) for variously different merger models. In \\S 4, we predict a possible relationship between MDFs of stellar haloes and physical properties of their host elliptical galaxies. In this section, we also discuss  the origin of M31's metal-rich stellar halo. We summarize our  conclusions in \\S 5. ",
        "conclusions": "We have numerically investigated MDFs of stellar haloes of elliptical galaxies formed by major merging in order to elucidate the origin of the observed MDF of NGC 5128. We mainly investigated the MDF of halo stars with $R$ $>R_{\\rm d}$ (where $R$ and $R_{\\rm d}$ are the radius from the center of a merger remnant and the initial size of the merger progenitor spiral) in the merger remnant for several models with different physical parameters of the progenitor spirals. Our numerical study clearly demonstrates the importance of {\\it dynamical processes of galaxy merging} (i.e., tidal stripping of metal-poor disc stars) in determining MDFs of stellar haloes of elliptical galaxies. We summarize our principal results as follows.  (1) The stellar halo of an elliptical galaxy formed by major merging is predominantly populated by moderately metal-rich stars with [m/H] $\\sim$ $-0.4$, which come from the outer parts of the discs of the merging spirals. Furthermore, the fraction of metal-poor stars with [m/H] $<$ $-1.0$ in the halo is very small ($\\sim$ 17 \\%). The MDF of the  merger remnant does not show a strong radial dependence for 0.5 $<$ $R$ $<$ 2.0 (in our units). These results are broadly consistent with recent {\\it HST} observations on MDFs of the stellar halo of NGC 5128.  (2) Irrespective of the physical parameters of merger progenitor discs, the MDFs of stellar haloes of merger remnants show both the dominant moderately metal-rich populations and the minor metal-poor ones. Thus future systematic  observations on MDFs of elliptical galaxies can assess the relative importance of major galaxy merging in the formation of elliptical galaxies by confirming the above generic trend of MDFs.  (3) The MDFs of elliptical galaxies formed by major merging depend  strongly on the initial MDFs of merger progenitor discs, because halo stars are dominated by those which are initially in the outer part of the discs and tidally stripped during merging.  (4) MDFs of stellar haloes do not depend so strongly on the MDFs of the progenitor spiral bulges, essentially because bulge stars are not be tidally stripped so efficiently during merging owing to bulge's compact configuration. However, the shapes of MDFs of a merger remnant at $-0.2$ $\\le$ [m/H] $\\le$ 0 can be affected by the bulge-to-disc-ratio of the merger progenitor spiral.  (5) The shape of the halo MDF of a merger remnant at [m/H] $\\le$ $-1.0$ depends on the mass fraction and the mean metallicity of the progenitor spiral's halo. Therefore the observed shape of the halo MDF at [m/H] $\\le$ $-1.0$ in NGC 5128 can give some constraints on the stellar populations of the progenitor spiral's stellar halo.  (6) The observed similarity in MDFs of stellar haloes between M31 and NGC 5128 can be naturally explained, if the M31 bulge was also formed by major merging between two spirals with the masses (a factor of $2-4$)  smaller than those that are assumed to merge to from NGC 5128 in the present study. The observed markable difference in MDFs between stellar haloes and GCs for NGC 5128 and M31 can be also explained by the merger scenario.   (7) Our simulations predict that if most elliptical galaxies are formed by major merging, the photometric properties (e.g., colors and magnitudes) of a stellar halo in an elliptical galaxy can correlate with those of the galaxy (i.e., redder Es have redder stellar haloes). Also our study predicts that there can be CM relations in stellar haloes of elliptical galaxies."
    },
    "0212/astro-ph0212290_arXiv.txt": {
        "abstract": "We present results from a detailed dynamical analysis of five high surface brightness, late type spiral galaxies NGC 3810, NGC 3893, NGC 4254, NGC 5676, and NGC 6643 which were studied with the aim to quantify the luminous-to-dark matter ratio inside their optical radii. The galaxies' stellar light distribution and gas kinematics have been observed and compared to hydrodynamic gas simulations which predict the gas dynamics arising in response to empirical gravitational potentials, which are combinations of differing stellar disk and dark halo contributions. The gravitational potential of the stellar disk was derived from near-infrared photometry, color-corrected to yield a constant stellar mass-to-light ratio (M/L); for the dark halo, the mass density distribution of an axisymmetric isothermal sphere with a core was chosen. Hydrodynamic gas simulations were performed for each galaxy for a sequence of five different mass fractions of the stellar disk and for a wide range of spiral pattern speeds. These two parameters mainly determine the modelled gas distribution and kinematics. The agreement between the simulated and observed gas kinematics permitted us to conclude that the galaxies with the highest rotation velocities tend to possess very massive stellar disks which dominate the gas dynamics within the optical radius. In less massive galaxies, with a maximal rotation velocity of $<$\\,200\\,km\\,s$^{-1}$, the mass of the dark halo at least equals the stellar mass within 2 -- 3 disk scale lengths. The maximal disk stellar mass-to-light ratio in the $K$-band was found to lie at about M/L$_K \\approx 0.6$. Furthermore, the gas dynamic simulations provide a powerful tool to accurately determine the dominant spiral pattern speed for galaxies, independent of a specific density wave theory. It was found that the location of the corotation resonance falls into a narrow range of around three exponential disk scale lengths for all galaxies from the sample. The corotation resonance encloses the strong part of the stellar spiral in all cases. Based on the experience gained from this project, the use of a color-correction to account for local stellar population differences is strongly encouraged when properties of galactic disks are studied that rely on their stellar mass distributions. ",
        "introduction": "Much evidence has accumulated in recent years that the stellar disks of high surface brightness spiral galaxies dominate the mass budget of the inner regions. Most of the studies arguing for a maximal disk scenario, are based on the detailed analysis of high resolution rotation curve measurements \\citep{bla99,pal00,rat00,sal01}. These studies, however, generally derive the rotational support of the stellar disk from an axisymmetric disk mass model and allow no consideration of non-circular rotation components. More sophisticated modelling strategies have been applied to spiral galaxies with variable significance \\citep{eri99,pig01}. For the most part those studies also support heavy disks for high surface brightness galaxies, but they also find candidates for which lower disk mass fractions are more likely. The most convincing direct arguments for the maximal disk scenario come from the studies of strong bars in spiral galaxies. These features induce a strong dynamic trace in the velocity field and provide a good laboratory for estimating the stellar mass component. Based on fluid dynamic modelling, a maximal disk solution is found by \\citet{eng99} for the Milky Way and by \\citet{wei01} for NGC 4123. Furthermore, theoretical considerations hint for the requirement of a non-massive halo contribution in the central regions of strongly barred galaxies with a very cold disk, because otherwise dynamical friction would slow down the bar very quickly, leading to its destruction \\citet{deb98,deb00}. On the other hand, for initially hotter disks, this effect is largely reduced \\citet{ath02} and references therein.   However, there is also evidence that even high surface brightness spiral galaxies might be dominated by the dark matter mass component in their central disk regions as suggested by high resolution simulations of cosmological dark matter halos \\citep[e.g.]{moo99a,ghi00,fuk01}. \\citet{bot93} inferred from stellar velocity dispersions in spiral galaxy disks that a more massive halo component is needed to explain the findings. There are recent studies that argue for lighter disk models by making use of previously hardly exploited relations. \\citet{CR99} applied a statistical Tully-Fisher relation analysis to a large sample of galaxies trying to relate the maximal rotation velocity of a galaxy to its disk size. \\citet{mal00}, and more recently \\citet{tro02} used the geometry of gravitationally lensed systems to disentangle the effects of the stellar disk and the halo masses. These groups found that the dark halo also dynamically plays an important role in the galaxies' inner regions.  In the present study, we quantify the relative mass fractions of the stellar disk and the inner dark halo by analyzing the non-axisymmetric force component in the gravitational force field that arises in response to the spiral structure of the observable stellar mass distribution. We address this issue by exploring the influence of the spiral arms on the gas velocity field that builds up in response to the total gravitational potential of the galaxy. We model the gravitational potential from a mass map of the stellar distribution and a dark halo model. With the use of hydrodynamical gas simulations, we predict the gas velocity fields for different empirical mass models of the sample galaxies, assembled from a variety of stellar disk and dark matter mass components. The two parameters that dominate the resulting forces in the disk are the stellar disk mass contribution and the pattern speed of the stellar spiral, $\\Omega_{\\rm p}$. Thus, in this study we explore the effects of those parameters and compare the simulated gas morphology and velocity wiggles with our observations, enabling the identification of the best fitting scenario for each individual galaxy. Comparable studies have been applied to barred spiral galaxies taking advantage of the strong dynamical trace of the bar in the velocity field \\citep{wei01}.  First results from our study have been presented already in \\citet{kra01}, (hereafter, KSR01) for NGC 4254 where it was proven that this strategy yields valuable information which allows us to constrain disk and halo fractions also for spiral galaxies without a strong bar. In the present paper we only briefly discuss the modelling and comparison strategies involved (Section~\\ref{modelling}) but instead focus on the conclusions of the analysis. More details about the modelling issues can be found in KSR01, \\citet{sly02b} and \\citet{kra02}. ",
        "conclusions": "We have presented a new, direct and independent estimate of the stellar mass fraction in a small sample of spiral galaxies. With respect to ratio of luminous and dark matter in their inner parts, the population of high surface brightness spiral galaxies does not seem to comprise an entirely homogeneous class of objects. Combining the maximum rotation velocity, presented in Table~\\ref{properties}, with the estimates for the galaxies' most probable fractions of the stellar disk mass, reveals an intriguing trend: the most massive of the analyzed galaxies, NGC 3893 and NGC 5676, tend to also possess the most massive stellar disks which dominate the dynamics of the central regions. The other galaxies from the sample are less massive systems that exhibit a maximum rotation velocity $v_c$\\,$<$\\,200\\,km\\,s$^{-1}$. For those objects the total mass of the dark halo within the optical radius is higher and was found to, at least, equal the stellar mass. This trend is graphically presented in Figure~\\ref{vc2fdcomp}.  Our results can be compared to \\citet{ath87}. These authors inferred disk and halo mass fractions from spiral structure constraints. The small points displayed in Figure~\\ref{vc2fdcomp} represent a subsample of the total sample (given in Athanassoula et al.~1987, Table A2), selected for late type galaxies with a bulge mass $M_{\\rm bulge} < 20$\\,\\% $M_{\\rm disk}$ and a maximum rotation velocity of more than 100\\,km\\,s$^{-1}$. For comparison the stated disk-to-halo mass ratio was converted into our scheme of quantifying the disk mass fraction ${\\rm f_d}$.  It can be clearly seen that the trend that emerged from our small dataset is supported by the results from \\citet{ath87}. In this sample there is no late type spiral galaxy with a maximum rotation velocity of over 200\\,km\\,s$^{-1}$ that seems to be in agreement with a substantially submaximal stellar disk. For less massive systems this is not the case anymore. Apparently galaxies with a wide range of dark halo and disk mass decompositions can exist and maintain spiral structure.  In light of this, the results from \\citet{CR99} might be explained that in a statistical sample of spiral galaxies the heavy and clearly maximum disk galaxies could account for only a relatively small fraction and most of the ``lighter'' high surface brightness galaxies exhibit already a considerable dark mass fraction. Indeed, in the Courteau \\& Rix sample, only roughly a third of the galaxies exhibit maximal rotation velocities $v_c$\\,$>$\\,200\\,km\\,s$^{-1}$.  However, beyond the discussion of the disk mass fraction in spiral galaxies, at present there is growing evidence that galaxy size dark halos are not cuspy in their centers, as claimed by standard cold dark matter (CDM) models. \\citep{moo94,flo94,moo99b,sal00,bor01,deb01,vdb01}. This circumstance was first found for low surface brightness and dwarf galaxies, that supposedly are comprised of a massive dark halo. Most of the authors argued that this also applies to high surface brightness galaxies. For the present analysis, the functional form of the dark halo profile has only a second order effect on the results as long as the models describe well the overall shape of the observed rotation curves. Thus, we decided not to distinguish between different dark halo profiles. However, the results from this study, which mostly confirm the strong dynamic influence of the stellar disk in massive high surface brightness spiral galaxies within the optical radius, provide very little room for a strongly cusped halo core."
    },
    "0212/astro-ph0212259_arXiv.txt": {
        "abstract": "We have found for the first time a Balmer edge feature in the Big Blue Bump emission of a quasar. The feature is seen in the polarized flux spectrum of the quasar, where all the emissions from outside the nucleus are scraped off and removed. The existence of the Balmer-edge absorption feature directly indicates that the Big Blue Bump is indeed thermal and optically-thick. ",
        "introduction": "Among the various components of a quasar emission, the optical/UV component, called Big Blue Bump (BBB), is energetically the most dominant. The BBB is often assumed to be from optically-thick thermal emission from an accretion flow, such as an accretion disk. However, this fundamental issue (i.e. the emission mechanism of the BBB) is actually very far from being solved, since observations are hardly said to be well described by disk models.  Among several serious problems are the continuum slope and apparent lack of continuum edges. In order to explain the observed optical/UV slope (e.g. $F_{\\nu} \\propto \\nu^{-0.3}$, Francis et al. 1991), the disk has to be rather cool.  However, naively, in a cool disk, its atmosphere would show large continuum edges (but see below). Apparently, we do not see such features. ",
        "conclusions": "Figure~1 shows our recent Keck spectropolarimetry data of a quasar. Thick line is the polarized flux spectrum, while the dotted line is the total flux spectrum scaled to roughly match the polarized flux at the red side. Emission lines are essentially all absent in the polarized flux (see Fig.2 which shows that the polarization clearly decreases at the line wavelengths).  The slope of the polarized flux at the red side is roughly the same as that of the total flux. However, we find that the slope changes at the bluer side of $\\sim$4000\\AA\\ in the rest frame, just where the 3000\\AA\\ bump starts in the total flux (there might also be a possible up-turn at the bluer side of $\\sim$3600\\AA).  This is an expected Balmer-edge absorption feature, and directly indicates that the Big Blue Bump is indeed thermal and optically thick, without involving any particular model of the nucleus.  \\bigskip   \\unitlength 1cm  \\begin{figure} \\hbox to \\textwidth{\\hfil \\begin{minipage}{7cm} \\begin{picture}(5,5) \\put(0,0){\\psfig{figure=kishimoto_fig1.ps,width=6cm}} \\end{picture} \\caption{{\\footnotesize The polarized flux and total flux of Ton202, taken at Keck in May 2002.}} \\end{minipage} \\hfil \\begin{minipage}{7cm} \\begin{picture}(5,5) \\put(0,0){\\psfig{figure=kishimoto_fig2.ps,width=6cm}} \\end{picture} \\caption{{\\footnotesize The normalized Stokes $q$ and $u$, with the scaled total flux at the bottom for reference.}} \\end{minipage} \\hfil} \\end{figure}"
    },
    "0212/astro-ph0212403_arXiv.txt": {
        "abstract": "We present results from the first hydrodynamical star formation calculation to demonstrate that close binary stellar systems (separations $\\lsim 10$ AU) need not be formed directly by fragmentation.  Instead, a high frequency of close binaries can be produced through a combination of dynamical interactions in unstable multiple systems and the orbital decay of initially wider binaries.  Orbital decay may occur due to gas accretion and/or the interaction of a binary with its circumbinary disc.  These three mechanisms avoid the problems associated with the fragmentation of optically-thick gas to form close systems directly. They also result in a preference for close binaries to have roughly equal-mass components because dynamical exchange interactions and the accretion of gas with high specific angular momentum drive mass ratios towards unity.  Furthermore, due to the importance of dynamical interactions, we find that stars with greater masses ought to have a higher frequency of close companions, and that many close binaries ought to have wide companions. These properties are in good agreement with the results of observational surveys. ",
        "introduction": "The process of star formation preferentially produces binary stellar systems (e.g.\\ Duquennoy \\& Mayor 1991).  The favoured mechanism for explaining this high frequency of binaries is the collapse and fragmentation of molecular cloud cores (e.g.\\ Boss \\& Bodenheimer 1979; Boss 1986; Bonnell et al.\\ 1991; Nelson \\& Papaloizou 1993; Burkert \\& Bodenheimer 1993; Bate, Bonnell \\& Price 1995). However, while fragmentation can readily create wide binary systems (separations $\\gsim 10$ AU), there are severe difficulties with fragmentation producing close binaries directly.  This is a significant deficiency since approximately 20\\% of solar-type stars have main-sequence companions that orbit closer than 10 AU (Duquennoy \\& Mayor 1991), and the frequency of massive spectroscopic binaries appears to be even higher (Garmany, Conti \\& Massey 1980; Abt et al.\\ 1990; Morrell \\& Levato 1991; Mason et al.\\ 1998).  As a molecular cloud core begins to collapse, the formation of wide binaries through fragmentation is possible because the gas easily radiates away the gavitational potential energy that is released.  The gas remains approximately isothermal and, thus, the Jeans mass decreases with density as $\\rho^{-1/2}$.  However, at densities of $\\gsim 10^{-13}$ g~cm$^{-3}$ or $n({\\rm H}_2)\\gsim 10^{10}$ cm$^{-3}$ (Larson 1969; Masunaga \\& Inutsuka 2000) the rate of heating from dynamical collapse exceeds the rate at which the gas can cool. The gas heats up, and the Jeans mass begins to increase so that a Jeans-unstable region of gas becomes Jeans-stable.  This results in the formation of a pressure-supported fragment with a mass of several Jupiter-masses and a radius of $\\approx 5$ AU (Larson 1969).  Fragmentation on smaller scales is inhibited by thermal pressure.  Therefore, initial binary separations must be $\\gsim 10$ AU.  The formation of such pressure-supported fragments is frequently refered to as the opacity limit for fragmentation (Low \\& Lynden-Bell 1976; Rees 1976) and may set a lower limit to the mass of brown dwarfs (Boss 1988; Bate, Bonnell \\& Bromm 2002).  A possibility for fragmentation at higher densities (hence on smaller length scales) exists when the pressure-supported fragment has accreted enough material for its central temperature to exceed 2000 K.  At this temperature, molecular hydrogen begins to dissociate, which provides a way for the release of gravitational energy to be absorbed without significantly increasing the temperature of the gas.  Thus, a nearly isothermal second collapse occurs within the fragment that ultimately results in the formation of a stellar core with radius $\\approx 1 R_\\odot$ \\cite{Larson1969}. Several studies have investigated the possibility that fragmentation during this second collapse forms close binary systems directly (Boss 1989; Bonnell \\& Bate 1994; Bate 1998, 2002).  Boss \\shortcite{Boss1989} found that fragmentation was possible during this second collapse, but that the fragments spiralled together due to gravitational torques and did not survive. Bonnell \\& Bate \\shortcite{BonBat1994} found that fragmentation to form close binaries and multiple systems could occur in a disc that forms around the stellar core.  However, both these studies began with somewhat arbitrary initial conditions for the pressure-supported fragment.  Bate \\shortcite{Bate1998} performed the first three-dimensional calculations to follow the collapse of a molecular cloud core through the formation of the pressure-supported fragment, the second collapse, and the formation of the stellar core and its surrounding disc. In these and subsequent calculations (Bate 2002), Bate found that the second collapse did not result in sub-fragmentation due to the high thermal pressure and angular momentum transport via gravitational torques.  In this paper, we present results from the first hydrodynamical star formation calculation to produce dozens of stars and brown dwarfs while simultaneously resolving beyond the opacity limit for fragmentation.  Despite the fact that no close binaries (separations $\\lsim 10$ AU) are formed by direct fragmentation, we find that the calculation eventually produces several close binary systems through a combination of dynamical interactions in multiple systems, and the orbital decay of wide binaries via gas accretion and their interactions with circumbinary discs.  The paper is structured as follows. In section 2, we briefly describe the numerical method and the initial conditions for our calculation.  In section 3, we present results from our calculation and compare them with observations.  Finally, in section 4, we give our conclusions. ",
        "conclusions": "We have presented results from a hydrodynamic calculation of the collapse and fragmentation of a turbulent molecular cloud to form 50 stars and brown dwarfs.  The calculation mimics the opacity limit for fragmentation by using a barotropic equation of state to model the heating of collapsing gas at high densities. This results in a minimum mass of $\\approx 10$ Jupiter masses for the lowest-mass brown dwarfs (see Bate et al.\\ 2002) and prevents fragmentation on scales smaller than $\\approx 10$ AU. Despite this lower limit on the initial minimum separation between fragments, we find that as the stellar groups and unstable multiple systems evolve in the gas-rich environment, a high frequency of close binary systems (separations $\\lsim 10$ AU) is produced.  Examining the history of these close binary systems, we find that they are formed through a combination of dynamical interactions in unstable multiple systems, and orbital decay due to accretion and/or the interaction of binary and triple systems with circumbinary and circumtriple discs.  These formation mechanisms allow realistic numbers of close binary systems to be produced without the need for fragmentation on length scales $<10$ AU.  This avoids the difficulties associated with the fragmentation of optically-thick gas during the collapse initiated by the dissociation of molecular hydrogen (Boss 1989; Bonnell \\& Bate 1994; Bate 1998, 2002).  As a consequence of the dependence of close binary formation on dynamical exchange interactions and the accretion of material with high specific angular momentum, we find that close binaries tend not to have extreme mass ratios.  All of our systems have mass ratios $q \\gsim 0.3$. Furthermore, the frequency of close binaries is dependent on mass in that massive stars are more likely to have close companions than lower mass stars. These properties are in good agreement with the results of observational surveys.  At the end of our calculation, many of the close binaries are members of hierarchical triple systems.  Although these systems may not yet have finished evolving, the implication is that many close binaries ought to have wider companions. Recent observations support this hypothesis, but larger surveys to determine the frequency of wide companions to close binary systems are necessary to demonstrate it conclusively."
    },
    "0212/astro-ph0212129_arXiv.txt": {
        "abstract": "We present a detailed analysis of the error budget for the TreePM method for doing cosmological N-Body simulations. It is shown that the choice of filter for splitting the inverse square force into short and long range components suggested in Bagla (2002) is close to optimum. We show that the error in the long range component of the force contributes very little to the total error in force. Errors introduced by the tree approximation for the short range force are different from those for the inverse square force, and these errors dominate the total error in force. We calculate the distribution function for error in force for clustered and unclustered particle distributions. This gives an idea of the error in realistic situations for different choices of parameters of the TreePM algorithm. We test the code by simulating a few power law models and checking for scale invariance. ",
        "introduction": "Cosmological N-Body simulations have played a crucial role in improving our understanding of formation of large scale structure. These have filled large gaps in a domain where analytical solutions do not exist.  N-Body simulations have also played a useful role by testing scaling relations derived from physically motivated ansatze.  Limitations of computing resources and our ability to simulate physical processes numerically have meant that N-Body simulations give only approximate solutions.  It is possible to test whether the approximations in evolution of the system affects estimation of physical quantities of interest so the results are generally more reliable than those based on approximate evolution of the system.  A large number of methods have been used for doing Cosmological N-Body simulations \\citep{edbert_araa}.  All of these use some approximation for calculation of force.  Approximations are used because direct summation over all pairs of particles scales as $O(N^2)$, where $N$ is the number of particles, making the calculation very time consuming for large $N$. The use of these approximations reduces the number of calculations required to $O(N\\ln{N})$ or less.  These approximations also introduce inaccuracies in the computed force.  The TreePM code \\citep{treepm}, that we study here, is a hybrid technique for carrying out large N-Body simulations to study formation and evolution of large scale structures in the universe. It is a combination of the \\citet{bh86} Tree code and a Particle-Mesh code \\citep{sim_book}. The TreePM method combines high resolution of tree codes with the ability of PM codes to compute the long range force with periodic boundary conditions.  In this paper we carry out a comprehensive study of the TreePM code. We analyse errors in estimation of force in both the tree and the PM components, and also study the distribution of errors for different distributions of particles.  We also study the variation of CPU time required for the TreePM code as we vary parameters describing the mathematical model of this method. ",
        "conclusions": "In this paper we have presented a detailed study of performance characteristics of the TreePM method.  We have analysed different sources of error and suggested remedies for the main source of errors.  The analysis of errors in realistic situations shows that the TreePM method performs very well and gives acceptably small errors.  This method compares favourably with other comparable methods such as implementations of the tree code like GADGET \\citep{gadget} and hybrid methods such as the TPM \\citep{tpm,tpm2}.  From the numbers available in these papers, we find that the errors in the recommended configuration of the TreePM method are comparable with those in GADGET and lower than those in the TPM method.  In terms of CPU time taken per step per particle, we again find that the numbers are comparable. Of course, it is not possible to make a detailed comparison of this quantity as the whole approach is different.  E.g., we do not use multiple time steps whereas GADGET relies heavily on these to optimise the speed.  On the other hand GADGET and TreePM have a uniform force resolution whereas TPM does not and hence the time taken is likely to vary more strongly with the amplitude of clustering for the TPM code as compared to the other two.  Splitting of force into a short range and a long range part has a useful implication in terms of parallel implementation.  Unlike the tree methods, the TreePM method involves a considerably smaller overhead in terms of inter-process communication and hence is likely to score over other methods in terms of scaling for a very large number of CPUs. In summary we can state that TreePM is a competitive method for doing Cosmological N-Body simulations.  Explicit use of three parameters gives users control over errors and CPU time required."
    },
    "0212/hep-th0212312_arXiv.txt": {
        "abstract": "The possibility of obtaining singularity free cosmological solutions in four dimensional effective actions motivated by string theory is investigated. In these effective actions, in addition to the Einstein-Hilbert term, the dilatonic and the axionic fields are also considered as well as terms coming from the Ramond-Ramond sector. A radiation fluid is coupled to the field equations, which appears as a consequence of the Maxwellian terms in the Ramond-Ramond sector. Singularity free bouncing solutions in which the dilaton is finite and strictly positive are obtained for models with flat or negative curvature spatial sections when the dilatonic coupling constant is such that $\\omega < - 3/2$, which may appear in the so called $F$ theory in 12 dimensions. These bouncing phases are smoothly connected to the radiation dominated expansion phase of the standard cosmological model, and the asymptotic pasts correspond to very large flat spacetimes. ",
        "introduction": "Superstring is the most promising candidate to describe a unified theory of all interactions, gravity included. There are five consistent superstring theories in 10 dimensions, which are connected among themselves through duality transformations. To each superstring theory, there is a corresponding supergravity theory in 10 dimensions. All of them can be obtained from the 11 dimensional supergravity theory. This indicates that those superstring theories are different manifestations of a unique 11 dimensional framework, that has been named $M$ theory~\\cite{polchinski,green,kiritsis}. Moreover, the superstring type-IIB can be recast in a more geometrical form in a 12 dimensional model, suggesting that perhaps a yet more fundamental framework may exist in 12 dimensions, which has been called $F$ theory~\\cite{pope}.  The physical properties of superstring theories become relevant at energy scales comparable with the Planck scale. This renders very improbable that superstring phenomenology may be tested in the near future in some laboratory experiment (see, however, Ref.~\\cite{randal} in which the Planck mass is lowered to TeV scale by accounting for large extra dimensions). According to the hot big bang scenario, however, energy scales even as high as the usual Planck scale ($\\mP\\sim 10^{19}$GeV) may have been reached in the very early universe. Hence, for the moment, cosmology seems to be the most natural arena where the consequences of superstring theories may be tested. The pre-big bang paradigm~\\cite{PBB} was one of the first ideas to implement superstring theories in this framework. Some relics of a cosmological string phase may also be identified~\\cite{bdp}, opening perhaps the possibility of testing superstring models. Furthermore, superstring theories open the possibility that some typical drawbacks of the standard cosmological model, such as the existence of an initial singularity, may be solved in the context of superstring cosmological models. The goal of the present paper is to show that, under certain conditions, it is possible to obtain completely regular bouncing cosmological models in the context of effective actions constructed from superstring theories (not involving, in particular, negative energies~\\cite{ppnpn2}), for which, moreover, the dilaton is strictly positive (nonvanishing) at all times and never diverges.  String cosmology is based on the low energy limit of string or superstring theories. In the most general case of the supersymmetric string theory, there are two sectors, related with the choice of periodic or antiperiodic boundary conditions on the spinor fields, namely the Ramond and Neveu-Schwarz (NS) sectors~\\cite{polchinski,green}. Since fermions can be either left or right moving, this leads to four possible combinations of these sectors. The bosonic fields arise both from the NS-NS and Ramond-Ramond (RR) sectors. The NS-NS sector provides the Einstein-Hilbert term, as well as a three-form, called the axionic field, and the dilaton. The latter is directly related with the string perturbative expansion parameter and takes the form of a Brans-Dicke-like scalar field, nonminimally coupled to both the Einstein-Hilbert and the axionic fields. In the RR sector, $p$ forms appear, which are minimally coupled to the dilaton field. The dimensions of the $p$ forms depend on the specific superstring theory which is under consideration. In cosmological applications, we are interested in scalar fields that emerge from these different $p$ forms. They differ by the way they couple with the dilatonic field and between themselves.  Out of the many different possibilities stemming from string theory, one can construct in general an effective action suitable for cosmological applications with two main features: scalar fields coming from the NS-NS sector and nonminimally coupled to the dilaton, and scalar fields from the RR sector which are minimally coupled to the dilaton. All the possibilities are not exhausted by these two frameworks, but they summarize the general aspects of what has been proposed in the literature as far as effective actions coming from string theory are concerned. One can also obtain phenomenological matter fields by averaging on some components of those original $p$ forms.  This brief description explicits the great richness of the string effective action procedure, which implies a large variety of possible cosmological models. Notice that these effective actions exhibit great similarities with those that can be obtained from multidimensional and supergravity theories. To select one cosmological model that could be a candidate to describe the physical world, two possible prescriptions are: either the cosmological model is completely regular, with no curvature or expansion parameter singularity, or it is compatible with observation; ultimately, both criteria should be satisfied. In the present work we will concentrate on the first. The second criterion, which presents some specific challenges, will be treated in the future~\\cite{future}.  The search for singularity free cosmology in string theories is not a new subject~\\cite{lidsey,vasquez,picco,kirill,branden}. The string action at tree level does not lead in general to singularity free cosmological solutions, at least when the strict string case ($\\omega = - 1$, $\\omega$ being the dilatonic coupling parameter) is considered. The pre-big bang model~\\cite{PBB}, which is an example of a string cosmology, requires the introduction of nonlinear curvature terms in order to achieve a smooth transition from a curvature growing phase to a curvature decreasing phase. If large negative values of the dilatonic coupling parameter $\\omega$ are allowed, it is possible, in some cases, to obtain completely regular models, including in the dilatonic sector~\\cite{kirill}. This may be achieved mainly in models with spatial sections with negative curvature.  Here, it will be shown that regular cosmological models may also be obtained if a radiation fluid is coupled to the string action at the tree level. Such a radiation fluid can have a fundamental motivation, for example, in the case of the superstring type IIB theory, where a 5-form appears in the RR sector. Truncation and dimensional reduction of this 5-form lead to a Maxwell term in four dimensions with the desired features~\\cite{fabris}. Hence, the model to be studied here is totally based on superstring theories.  The string motivated phenomenological term included under the form of a radiation fluid makes it possible to connect smoothly such string cosmological models to the radiation phase of the standard cosmological model before nucleosynthesis.  In Ref.~\\cite{picco}, models motivated by string theory similarly including a radiation fluid have been studied, restricted to flat spatial sections and $\\omega > -3/2$.  In such cases, bouncing solutions have been obtained only for $\\omega < -4/3$. Furthermore, for these solutions, the dilaton vanishes in the infinite past, raising doubts on the validity of the tree level action in such a region. In the present paper, the curvature of the spatial sections and the value of $\\omega$ are kept arbitrary. New bouncing regular solutions are then obtained, for which, as mentioned above, the dilaton remains finite and nonvanishing at all times. When $\\omega > -3/2$, which includes the strict string case ($\\omega = -1$), the solutions can only be bouncing provided the spatial sections have negative curvature and if the dilaton is always negative, which is not consistent with the higher dimensional framework of stringlike theories, and implies a repulsive gravity.  We shall henceforth disregard such solutions. When $\\omega < -3/2$, the bouncing solutions are obtained for models with flat or negative curvature, and the dilaton is strictly positive at all times.  In the following section, we derive effective string motivated actions in four dimensions. In particular, we show how to obtain an effective string action in four dimensions with $\\omega < -3/2$ in the context of the so-called $F$ theory in twelve dimensions. In Sec. III we derive nonsingular cosmological solutions from these effective theories, which are thoroughly discussed in Sec. IV from the point of view of violation of energy conditions. We end up with the conclusions in Sec. V. ",
        "conclusions": "We have constructed fully regular cosmological solutions in the framework of effective actions derived from string theory principles. These solutions present bouncing behaviors for a wide range of parameters ($\\omega < -3/2$), and are singularity free; furthermore, the spacetimes they lead to are geodesically complete, thereby improving the so-called horizon problem of standard cosmology. Stemming from string theory in the context of the so-called $F$ theory in twelve dimensions, where it is possible to have $\\omega < -3/2$, they have a reasonably sound basis as long as the dilaton is strictly positive and finite in such a case. As a consequence, it is not necessary to go beyond the tree level approximation in any part of their histories: the analytic solutions exhibited above can describe the whole history of the cosmological models they represent. Their consequences may, in turn, be used as cosmological tests.  Remembering that the radiation fluid included here also has a motivation in the superstring type IIB action, this turns out to be, to our knowledge, the first case where a complete regular bouncing cosmological solution is obtained in the string framework and related theories, which moreover is smoothly connected with the standard cosmological model radiation dominated phase. This solution may have flat or negative curvature spatial sections.  In the axionic and RR cases with $k=0$ or $k=-1$, there are nonsingular bouncing solutions for $-3/2<\\omega<-4/3$ but with vanishing dilaton in the beginning, where the tree level action cannot be trusted, and bouncing solutions with an initial curvature singularity if $-4/3<\\omega<0$. If $k = - 1$ and $-3/2<\\omega<0$ (normal, including the pure string case), one can have singularity free bouncing solutions with, however, a negative definite dilaton field.  As all the models with $\\omega<-3/2$ have the interesting feature to approach flat spacetime in the infinite past (either in Milne coordinates for $k=-1$, or the infinitely large radiation dominated standard model with $k=0$), there is the possibility to implement a quantum spectrum of perturbations in the initial asymptotics without any trans-Planckian problem~\\cite{transP}, and, at the same time, to accomplish a smooth transition to the standard cosmological model when, after the bounce, a standard radiation dominated phase is recovered (asymptotically in the $k=0$ case), preserving some of its main achievements like, \\eg primordial nucleosynthesis. The bouncing solutions with $\\omega>-3/2$ and positive dilaton still present some sort of trans-Planckian problem as long as the string expansion parameter diverges initially and one must go beyond the tree level action in such cases.  Notice that, in the cases where the dilaton is strictly positive, the initial value of the dilatonic field can be made smaller than its final value. Hence, the gravitational coupling can be initially given a much greater value than it would have today. This opens the possibility to solve the hierarchical problem of the gravitational coupling, in a spirit similar to the so-called brane cosmology~\\cite{langlois}."
    },
    "0212/astro-ph0212513_arXiv.txt": {
        "abstract": "The variability of gamma-ray burst (GRB) is thought to be correlated with its absolute peak luminosity, and this relation had been used to derive an estimate of the redshifts of GRBs. Recently Amati et al. presented the results of spectral and energetic properties of several GRBs with known redshifts. Here we analyse the properties of two group GRBs, one group with known redshift from afterglow observation, and another group with redshift derived from the luminosity - variability relation. We study the redshift dependence of various GRBs features in their cosmological rest frames, including the burst duration, the isotropic luminosity and radiated energy, and the peak energy $E_p$ of $\\nu F_\\nu$ spectra. We find that, for these two group GRBs, their properties are all redshift dependent, i.e. their intrinsic duration, luminosity, radiated energy and peak energy $E_p$, are all correlated with the redshift, which means that there are cosmological evolution effects on gamma-ray bursts features, and this can provide an interesting clue to the nature of GRBs. If this is true, then the results also imply that the redshift derived from the luminosity - variability relation may be reliable. ",
        "introduction": "The study of gamma-ray bursts (GRBs) afterglows has enable the measurement of their distances, so far GRBs are known as an explosive phenomenon occurring at cosmological distances, emitting large amount of energy mostly in the gamma-ray range (see, e.g.  Piran 1999; Cheng \\& Lu 2001 for a review). So GRBs can provide useful information about the early epochs in the history of the universe.  Although a great achievements have been made about the GRB afterglows, we still know little about the origin of gamma-ray bursts, the reason is that the GRBs with known redshifts are relatively rare, now there are only about 20 GRBs with known redshifts, so it is difficult to do some statistics about GRBs features, such as their luminosity function, duration distribution, etc.. However, two important correlations have been discovered, i.e. between the degree of variability of the gamma-ray burst light curve and the GRB luminosity (Ramirez-Ruiz \\& Fenimore 1999; Feminore \\& Ramirez-Ruiz 2001), and between the differential time lags for the arrival of burst pulses at different energies and the GRB luminosity (Norris, Marani \\& Bonnell 2000), although these correlations are still tentative, they offer the possibility to derive independent estimates of the redshifts of GRBs.  Recently Amati et al. (2002) have reported the spectral and energetic properties of several GRBs with known redshifts, these bursts were all detected by BeppoSAX and have good-quality time-integrated spectra. In addition, Lloyd-Ronning \\& Ramirez-Ruiz (2002) have found that bursts with highly variable light curves have greater $\\nu F_\\nu$ spectral peak energies in their cosmological rest frames. These results reinforce the validity of the redshift estimates derived from the luminosity - variability relation. Here we will discuss the properties of two group GRBs, one group includes the bursts with known redshifts and well-defined spectra detected by BeppoSAX, and another group consists of bursts whose redshifts are derived from the luminosity - variability relation. We will show that the properties of these two group GRBs are all correlated with redshift, which suggests that the luminosity - variability relation may be reliable, and furthermore the GRBs' features are redshift dependent. ",
        "conclusions": "In previous section we have shown that the properties of our two group GRBs, one group includes 9 GRBs with secure redshifts derived from the afterglow observation,  another group consists of 159 GRBs with known peak energy and their redshifts are derived from the luminosity - variability relation, are all correlated with the redshifts, which reinforces the validity of the luminosity - variability distance indicator. Our results further support the conclusion obtained by Lloyd-Ronning \\& Ramirez-Ruiz (2002), they found that there is positive correlation between the peak energy in the cosmological rest frame and the variability for gamma-ray bursts whose redshifts are derived either from optical spectral features or from the luminosity - variability distance indicator.  Since the BATSE GRB sample is flux truncated, i.e. only those bursts with flux exceeding the threshold flux can be detected, so it is necessary to discuss this selection effects on the statistical results. Here we use the Monte Carlo simulation method. We find that, after correcting the selection effects, the cosmological evolution trends are still exist, although they are much shallower than the trends found in the face values of the data. For example, if ignored the flux truncation effect, we obtained the relation $L \\propto (1+z)^{2.5\\pm 0.1}$, while when the selection effects are taken into account, we got the relation $L \\propto (1+z)^{1.7\\pm 0.5}$, which is consistent with the value $L \\propto (1+z)^{1.4\\pm 0.5}$ obtained by Lloyd-Ronning et al. (2002).  Fig.5 illustrates our simulated results, where the circles are the 300 bursts produced by the simulation, the solid line is the flux threshold for BATSE, and the dashed line is the flux threshold for Swift. From this figure it is obvious that the property of the \"observed\" sample should depend on the adopted flux threshold, for different flux threshold their statistical properties are different. For example, when taking the BATSE flux threshold, there are 100 bursts with flux lower than the threshold, while when taking the Swift flux threshold, there are only 7 bursts with flux lower than the threshold, so in this case the observed $L-z$ relation should close to the true $L-z$ relation. Therefore we expect that the luminosity - redshift relation for bursts observed by Swift should be shallower than that for bursts observed by BATSE.  Fig.6 and Fig.7 show that there is a good correlation between the peak energy and luminosity for both group GRBs, and the relation $E_p \\propto L^{1/2}$ can account for the observed data quite well. Up to now the location of the GRB emission site is still unsettled, although the internal shock model is thought to be more reasonable than the external shock model. Zhang \\& Meszaros (2002b) analysed various fireball models within a unified picture and investigated the $E_p$ predictions of different models. It is known that for internal shock model, if the GRB bulk Lorentz factors are not correlated with the luminosities, then there is the relation $E_p \\propto L^{1/2}$, while for external shock model $E_p \\propto \\Gamma ^4$, where $\\Gamma$ is the shock Lorentz factor. So our results suggest that the gamma-ray burst emission are more likely from the internal shock.  Frail et al. have discussed the afterglow properties of several GRBs with known redshifts, they assumed that the breaks in the afterglow light curves are caused by the sideways expansion of the jet, and then they concluded that the GRB emission energy is nearly a constant, $E\\sim 5\\times 10^{50}$ ergs (Frail et al. 2001). However, Fig.3 shows that the isotropic radiated energy increases with redshift, so if the conclusion of Frail et al. is true, then we notice that the jet opening angle must decrease with the redshift, as shown by Fig.8, where the data points are taken from the paper of Frail et al. (2001). This point is very interesting, since it can put constraints on the central engines of GRBs. Of course, this phenomena can also be explained within the framework of a structured universal jet model (Zhang \\& Meszaros 2002a; Rossi et al. 2002). In this model, an observer closer to the jet axis would detect a higher luminosity, thus at higher redshift, smaller viewing angle detections are preferred due to luminosity selection effect.  \\begin{figure} \\centerline{\\psfig{file=fig8.eps,width=0.6\\textwidth}} \\caption{The relation between the jet opening angle and redshift. The data are taken from the paper of Frail et al. (2001). } \\end{figure}  It should be noted that in our analysis we have not considered the errors coming from the luminosity - variability relation. We know that the relation between luminosity and variability is somewhat scatter, the correlation coefficient $n$ ($L\\propto V^n$) can be range from 2.2 to 5.8 (the best value is 3.3, see Fenimore \\& Ramirez-Ruiz 2001), so it is natural that the redshifts and luminosities inferred from this luminosity - variability relation should have large errors, and these errors should be somehow transferred to the final errors in the correlation indices. However this effect is very complicated, we hope that this effect can be taken into account in the future work.  In summary, in this paper we discuss the properties of two group GRBs, one group with known redshift from afterglow observation, and another group with redshift derived from the luminosity - variability relation. We find that the properties of these two group GRBs are all correlated with the redshifts, which reinforces the validity of the redshift estimates derived from the luminosity - variability relation. If this is true, then we see that the burst features, such as their intrinsic duration, luminosity, radiated energy and peak energy $E_p$, are all redshift dependent, which means that there are cosmological evolution effects on gamma-ray bursts features, and this can provide an interesting clue to the nature of GRBs."
    },
    "0212/astro-ph0212039_arXiv.txt": {
        "abstract": "We present the first results obtained at CFHT with the TRIDENT infrared camera, dedicated to the detection of faint companions close to bright nearby stars.  The camera's main feature is the acquisition of three simultaneous images in three wavelengths (simultaneous differential imaging) across the methane absorption bandhead at 1.6~$\\mu $m, that enables a precise subtraction of the primary star PSF while keeping the companion signal. The main limitation is non-common path aberrations between the three optical paths that slightly decorrelate the PSFs. Two types of PSF calibrations are combined with the differential simultaneous imaging technique to further attenuate the PSF: reference star subtraction and instrument rotation to smooth aberrations. It is shown that a faint companion with a $\\Delta H$ of 10 magnitudes would be detected at $0.5^{\\prime \\prime}$ from the primary. ",
        "introduction": "The next generation high order adaptive optics (NGAO) systems that are currently in development will produce high Strehl ratio images in the infrared on 10~m class telescopes. If those NGAO systems are to be used to search for very faint companions (brown dwarfs and exoplanets), it is of great importance to understand the limits for high contrast imaging on current adaptive optics (AO) systems and see how they would affect NGAO systems performances. Static aberrations have long been known to be the limiting factor for space-based telescopes when attempting high contrast imaging for faint companion detection (\\cite{Brown90}). If the telescope physical conditions change in time (like point spread function (PSF) breathing for HST), the structure becomes quasi-static and produces PSF structures that evolve slowly in time, making difficult a good PSF calibration when using reference stars. Long AO corrected exposures have shown that ground-based imaging is not free from this problem. When using an AO system that delivers partially diffracted PSFs, the final wavefront is an interference between the AO wavefront residuals and the optical path wavefront aberrations to the detector. If enough random aberrations are removed by the AO system, a coherent PSF starts to appear. This is also true for the quasi-static portion of the wavefront. As more turbulence is removed, the quasi-static wavefront errors not corrected by the AO system will introduce coherent structures that will not average out in time and will mask faint companions. A partially-corrected long exposure AO observation thus suffers from the same problem that limits space telescopes. The current limit for detecting faint companions with existing AO systems is not the atmospheric turbulence correction efficiency but rather a quasi-static PSF calibration problem. \\vspace{0.3cm}  In the past few years, our group has developed a specialized infrared camera, TRIDENT (\\cite{Marois2000a,Marois2002}), to overcome the PSF calibration problem. The main feature of TRIDENT is to acquire three simultaneous images in three distinct narrow spectral bands. It is possible to enhance the star/companion contrast after image combinations by selecting special spectral features that are typical of the companion and not of the star (\\cite{Smith87,Rosen87,Racine99,Marois2000b}). In TRIDENT, the three wavelengths (1.567~$\\mu$m, 1.625~$\\mu$m and  1.680~$\\mu$m, 1\\% bandwidth) have been selected across the 1.6 $\\mu$m methane absorption bandhead that is only present in the spectrum of cold ($T_{\\rm{eff}}$ $<$ 1470~K, \\cite{Feg96}) substellar objects. The 1.567~$\\mu$m image ($I_{\\lambda_1}$) shows the star and companion while the 1.625 and 1.680~$\\mu$m images ($I_{\\lambda_2}$ and $I_{\\lambda_3}$) show the star with a fainter companion due to the methane absorption bandhead. The 1.625 and 1.680~$\\mu$m images can thus be used as reference PSFs to correct the first and second order PSF structure evolution with wavelength. Since the three wavelengths are simultaneous, each observed wavelength sees the same PSF atmospheric distortion, a good PSF correlation between the three wavelengths can thus be achieved. This concept would remove atmospheric speckles and problems associated with the PSF evolution with telescope pointing, thermal changes and atmospheric r$_0$ variations. PSF subtraction is achieved by calculating simple and double image differences (SD and DD) obtained with the following algorithms (\\cite{Marois2000b}):  \\begin{center} $sd_{1-2} = I_{\\lambda_1} - I_{\\lambda_2}$ and $dd = sd_{1-2} - $k$*sd_{1-3}$ \\end{center}  The DD can bring the atmospheric speckle noise below the level of the photon noise and remove the PSF quasi-static structure to its second order evolution with wavelength. \\clearpage ",
        "conclusions": "Quasi-static aberrations in AO corrected images are the main problem when attempting high contrast imaging for faint companion searches. If quasi-static aberrations cannot be minimized or removed in NGAO systems, an optimized camera will be essential to calibrate and subtract them. If a multiple optical path camera is considered, non-common path aberrations need to be minimized to avoid PSF decorrelation. Reference star subtraction and aberration smoothing by instrument rotation improve by one order of magnitude the faint companion detection threshold. Typical attenuation in good seeing conditions are 10 magnitudes in $H$ band at $0.5^{\\prime \\prime}$ on a 3.6~m telescope.  \\vspace{0.5cm} This work is supported in part through grants from NSERC, Canada and from Fonds FQRNT, Qu\\'{e}bec."
    },
    "0212/astro-ph0212455_arXiv.txt": {
        "abstract": "A comprehensive study on compactness has been carried out on the 2dF Galaxy Group Catalogue constructed by Merch\\'an \\& Zandivarez. The compactness indexes defined in this work take into account different geometrical constraints in order to explore a wide range of possibilities. Our results show that there is no clear distinction between groups with high and low level of compactness when considering particular properties as the radial velocity dispersion, the relative fraction of galaxies per spectral type and luminosity functions of their galaxy members.  Studying the trend of the fraction of galaxies per spectral type as a function of the dimensionless crossing time some signs of dynamical evolution are observed. From the comparison with previous works on compactness we realize that special care should be taken into account for some compactness criteria definitions in order to avoid possible biases in the identification. ",
        "introduction": "Compact groups (CGs) are small systems of a few galaxies which are in close proximity one another. Their are excellent laboratories for studying galaxy interactions and, in particular, merging processes. Given their high galaxy density (equivalent to those at the centers of rich clusters) and small velocity dispersion ($\\sim 200 \\kms$), CG members are expected to finally merge into one giant elliptical galaxy within a few short crossing times.  Historically, CGs were of interest because of the obvious distortion of many of their member galaxies. The first systematic seek for CGs was pioneered by Rose (1977), using a surface number density contrast procedure. The most widely analysed samples are the Hickson Compact Groups (HCGs) (Hickson 1982, Hickson 1997), which have been selected on the basis of population, isolation (avoiding cores of rich clusters) and compactness. Their compactness criteria involve the computation of the mean surface brightness of a group which should be lesser than a maximun limit. This mean was calculated distributing the flux of the member galaxies over the circular area containing their geometric centers. Previous samples have been visually selected, and thus reflect some of the systematic biases intrinsic to bidimensional identification of systems. A geometric bias arise because prolate systems along the line of sight will be preferentially selected. A kinematic bias could enhance the selection of systems which are in a transient compact configuration due to galaxy internal motions. Mamon (1986) has suggested that about half of HCGs are superpositions of galaxies within loose groups (hereafter LGs). This suggestion could imply that groups properties in any particular sample may be strongly influenced by the criteria used to define the sample. Another important clue in order to test the compact groups environment is the color of their galaxy members. It is well known that elliptical galaxies recently formed from mergers of spiral galaxies should be bluer than normal elliptical galaxies. This kind of interactions should be more frecuently observed in a compact group environment. Nervertheless, a study on galaxy members of HCGs made by Zepf, Whitmore \\& Levinson (1991) have show that most of the early-type galaxies have optical colors indistinguishable form those of elliptical galaxies in other environments. Furthermore, there is evidence of a strange absence of strong signs of interactions, strong radio sources or far infrared radiation emission, etc. (Menon 1995, Pildis, Bregman \\& Schombert 1995, Sulentic 1997). Recently, Tovmassian, Yam \\& Tiersch (2001) and Tovmassian (2001) presents new evidence that indicates that almost all HCGs are dynamically associated generally with elongated LGs which are distributed along the elongation of the corresponding groups, suggesting that the HCGs are the compact cores of the latter. Consequently, an important question about the real nature of CGs arise: are CGs a distinct class by themselves or are extreme examples of systems having a particular range of galaxy density and population.  This question can be addressed studying the spatial distribution and environment of CGs. However, estimating the velocity dispersion and physical separations of galaxies in groups with a small number of galaxies is highly uncertain. Meaningful conclusions on dynamical properties about systems containing only four or five galaxies requires statistical analysis of large homogeneous samples.  \\begin{figure} \\epsfxsize=0.5\\textwidth \\hspace*{-0.5cm} \\centerline{\\epsffile{figzzz.ps}} \\caption{ The median mean nearest neighbour separation $\\langle R_{nn}\\rangle$ (upper panel) and the median virial radius $\\langle R_{vir}\\rangle$ (lower panel) as a function of redshift (filled circles) superimposed to the real data (dots) distribution. } \\label{fig0} \\end{figure}  \\begin{figure} \\epsfxsize=0.5\\textwidth \\hspace*{-0.5cm} \\centerline{\\epsffile{graf1.ps}} \\caption{ Example (a): Possible configuration of binary galaxy members of a group separated one another by large distances. Example (b): Configuration of the group galaxy members showing the possibility to find a loose group with a central concentration of galaxies. } \\label{fig00} \\end{figure}  \\begin{figure} \\epsfxsize=0.5\\textwidth \\hspace*{-0.5cm} \\centerline{\\epsffile{fig1.ps}} \\caption{ Upper left panel: Scatter plot of the normalized mean nearest neighbour separation of group members in the 2dFGGC versus the normalized virial radius. Upper right panel: Distribution of the normalized mean nearest neighbour separation. Lower right panel: Distribution of the normalized virial radius. Lower left panel: Distribution of the compactness index 1 ($CI_1$) defined in section 3 for groups in the 2dFGGC. } \\label{fig1} \\end{figure}  In order to overcome these biases, it has recently become feasible to find CGs using automatic identification. Such a procedure has the advantage of generate a sample that is homogeneous and complete within the criteria specified for the search. Barton et al. (1996) have used a selection criteria (friends-of-friends algorithm) based only on physical extent and association in redshift space. Eventhough Hickson's isolation criteria is not present at all in Barton's work, this technique is more effective in finding groups in regions of higher galaxy density because foreground and background galaxies are automatically eliminated by the velocity selection criteria. Other automatic algorithm for the selection of CGs from large galaxy catalogues has been developed by Iovino et al. (1999). The algorithm is such as to maximize the probability that the groups selected are physically related and partially reproduces the criteria used in the visual search by Hickson (1982), where his isolation criteria is slightly relaxed.\\\\ Since it is very difficult to identify compact groups at high redshifts, more reliable results can be obtained restricting the analysis to low redshift samples. Under this constraint applied on the Updated Zwicky Catalogue, Focardi \\& Kelm (2002) have shown that triplets are characterized by different properties than that obtained for higher order compact groups suggesting the existence of two different galaxy systems in the compact group samples. It is therefore of great interest to obtain larger and deeper samples of CGs, in order to put the CGs properties on a statistical basis. This will help to work out the controversy around the properties of CGs, contradictions that may only be apparent, given our still limited knowledge of the nature of these structures.  Currently, one of the largest group catalogues (hereafter 2dFGGC) was constructed by Merch\\'an \\& Zandivarez (2002). They have identified groups in the 2dF public 100K data release using a modified Huchra \\& Geller (1982) group finding algorithm that takes into account 2dF magnitude limit and redshift completeness masks. This catalogue constitutes a large and suitable sample for the study of both, processes in group environments and the properties of the group population itself. The global effects of group environment on star formation was analysed by Mart\\'{\\i}nez et al (2002a) using this catalogue. Dom\\'{\\i}nguez et al (2002) presented hints toward understanding local environment effects on the spectral types of galaxies in groups by studying the relative fractions of different spectral types as a function of the projected local galaxy density and the group-centric distance. Recently, an extensive statistical analysis on galaxy luminosity function in groups was carried out by Mart\\'{\\i}nez et al (2002b).  The aim of this work is to perform an analysis on the groups in the 2dFGGC by defining new compactness indexes which are assigned to every group in the sample. Several studies have been made using galaxy spectral type, velocity dispersion, luminosity and crossing time as a function of the compactness indexes. The outline of this paper is as follows. In section 2, we present the 2dF Galaxy Group Catalogue (2dFGGC) used throughout this work. Section 3 describes the compactness indexes definitions while in section 4 we analyse the possible dependence of our indexes with group and galaxy properties. Finally, in section 5 we summarize our conclusions. ",
        "conclusions": "In this work we report a statistical compactness analysis using the catalogue of galaxy groups constructed by Merch\\'an \\& Zandivarez (2002). Given the size of the original sample and 3-dimensional information our study intend to characterize groups of galaxies according to the level of compactness and analyse its influence on groups and galaxy members. For this purpose, we define two new compactness indexes based on geometrical criteria. Whereas index $CI_1$ prioritizes the distance to the nearest neighbour and the size of the system (eq. 1), index $CI_2$ enhances a possible core concentration in the system (eq. 2). Special cares have been taken over the construction of these indexes in order to avoid possible dependences on redshift and redshift completeness of the parent catalogue (Figures \\ref{fig2} and \\ref{fig3}).  With this characterization, we develop a wide analysis over many physical properties of groups and galaxy members. First, we observe that the compactness indexes distribution shows identical behaviours when the group sample is split for low ($\\sigma_r < 300 \\kms$) and high ($\\sigma_r \\ge 300 \\kms$) velocity dispersion. Furthermore, we observe that groups with high and low levels of compactness show the same normalized radial velocity dispersion distributions with a mean of $\\sim 200 \\kms$. This result is consistent with previous ones which reflect that the mean velocity dispersion of compact groups is quite similar to that found for loose groups (Hickson et al. 1992). Analysing the dependence of group compactness with numerical richness we observe that groups with high level of compactness have typically a low number of galaxy members, while most of the loose groups are characterized by a larger number of galaxy members.  On the other hand, another study has been made about the fraction of galaxies per spectral type and luminosity as a function of the two compactness indexes (Figures \\ref{fig8}, \\ref{fig9} and \\ref{fig10}). Our results do not show any particular correlation between the above parameters and the compactness level for groups in the sample.  The similar behaviour observed in groups with high level of compactness and loose ones probably suggests that this distinction is not fundamental. This result is supported by previous works which state an indistinguishable behaviour between compact and loose groups showing that many compact groups are located within overdense environments (de Carvalho et al. 1997, Barton et al. 1998, Zabludoff \\& Mulchaey 1998). Furthermore, an analysis on X-ray properties of groups shows that it is impossible to separate loose and compact groups on the luminosity-temperature relation, the luminosity-velocity dispersion relation or in the velocity dispersion-temperature relation stating that a more useful distinction is that between X-ray bright and X-ray faint systems (Helsdon \\& Ponman 2000).  The mean dimensionless crossing times obtained for a sample with high level of compactness is shifted toward higher values when comparing to the obtained for a sample of compact groups constructed by Hickson et al. (1992). This shift could be due to some biases in the compact group identification criteria. These biases could imply the detection of the cores of larger systems generating smaller dimensionless crossing times determinations.  The last correlation we studied is the fraction of galaxies per spectral type as a function of the dimensionless crossing time. Eventhough the correlations we found are not significant, it is worth to mention that for Type 2 galaxies, smaller is the fraction when higher is the dimensionless crossing time and the opposite trend is maintained for Type 4 galaxies. This latter result is consistent with the stated by Hickson et al. (1992) about groups with smaller crossing times typically containing fewer late-type galaxies."
    },
    "0212/astro-ph0212380_arXiv.txt": {
        "abstract": "We present results from the largest numerical simulation of star formation to resolve the fragmentation process down to the opacity limit.  The simulation follows the collapse and fragmentation of a large-scale turbulent molecular cloud to form a stellar cluster and, simultaneously, the formation of circumstellar discs and binary stars.  This large range of scales enables us to predict a wide variety of stellar properties for comparison with observations.  The calculation clearly demonstrates that star formation is a highly-dynamic and chaotic process.  Star formation occurs in localised bursts within the cloud via the fragmentation both of dense molecular cloud cores and of massive circumstellar discs. Star-disc encounters form binaries and truncate discs. Stellar encounters disrupt bound multiple systems. We find that the observed statistical properties of stars are a natural consequence of star formation in such a dynamical environment. The cloud produces roughly equal numbers of stars and brown dwarfs, with masses down to the opacity limit for fragmentation ($\\approx 5$ Jupiter masses). The initial mass function is consistent with a Salpeter slope ($\\Gamma=-1.35$) above 0.5 M$_\\odot$, a roughly flat distribution ($\\Gamma=0$) in the range $0.006-0.5$ M$_\\odot$, and a sharp cutoff below $\\approx 0.005$ M$_\\odot$. This is consistent with recent observational surveys. The brown dwarfs form by the dynamical ejection of low-mass fragments from dynamically unstable multiple systems before the fragments have been able to accrete to stellar masses. Close binary systems (with separations $\\lsim 10$ AU) are not formed by fragmentation in situ.  Rather, they are produced by hardening of initially wider multiple systems through a combination of dynamical encounters, gas accretion, and/or the interaction with circumbinary and circumtriple discs. Finally, we find that the majority of circumstellar discs have radii less than 20 AU due to truncation in dynamical encounters. This is consistent with observations of the Orion Trapezium Cluster and implies that most stars and brown dwarfs do not form large planetary systems. ",
        "introduction": "The collapse and fragmentation of molecular cloud cores to form bound multiple stellar systems has been the subject of many numerical studies (e.g.\\ Boss \\& Bodenheimer 1979; Boss 1986; Bonnell et al.\\ 1991; Nelson \\& Papaloizou 1993; Bonnell 1994; Burkert \\& Bodenheimer 1993; Bate, Bonnell \\& Price 1995; Truelove et al.\\ 1998). These calculations have resulted in the adoption of fragmentation as the favoured mechanism for the formation of binary and multiple stars, since it can produce a wide range of binary properties through simple variations of the pre-collapse initial conditions.  However, while individual binary systems can be reproduced by such fragmentation calculations, it is extremely difficult to use these calculations to predict the statistical properties of the stellar systems that should result from the fragmentation model. Quantities that we may wish to determine include the initial mass function (IMF), the relative frequencies of single, binary and multiple stars, the properties of multiple stars, the properties of circumstellar discs, and the efficiency of star formation. In order to predict these statistical properties, we need to produce a large sample of stars.  There are two possibilities.  We could perform many calculations of isolated cloud cores using a representative sample of initial conditions.  However, this has two disadvantages.  First, the conditions in molecular clouds are not sufficiently well understood to be able to select a representative sample of cloud cores for the initial conditions.  Second, the production of isolated stellar systems neglects interactions between systems that may be important in determining stellar properties, especially in young star clusters.  Examples of such interactions include binary formation via star-disc capture (Larson 1990; Clarke \\& Pringle 1991a,b; Heller 1995; McDonald \\& Clarke 1995; Hall, Clarke \\& Pringle 1996), truncation of protostellar discs (Heller 1995; Hall 1997), and competitive accretion leading to a range of stellar masses (Larson 1978; Zinnecker 1982; Bonnell et al. 1997, 2001a,b).  The second possibility is to perform a calculation of the collapse and fragmentation of a large-scale molecular cloud to form many stars simultaneously.  This is the approach taken in this paper.  Interactions between stars are automatically allowed for.  We must still specify global initial conditions, but the formation of individual cores within the cloud occurs self-consistently; we do not have to select initial conditions for each core arbitrarily.  The only disadvantage is that such a calculation is extremely computationally intensive.  Several such global calculations have been performed in the past. The earliest was that of Chapman et al.\\ \\shortcite{Chapmanetal1992} who followed the collapse and fragmentation of a shock-compressed layer of molecular gas between two colliding clouds.  The calculation produced many single, binary, and multiple protostars, but there was no attempt to derive the statistical properties of these systems. Klessen, Burkert \\& Bate \\shortcite{KleBurBat1998} followed the collapse of a large-scale clumpy molecular cloud to form $\\sim 60$ protostars.  They found that the mass function of the protostars could be fit by a lognormal mass function that has a similar width to the observed stellar initial mass function.  The protostellar masses were set by a combination of the initial density structure, competitive accretion, and dynamical interactions.  Further calculations in which the global initial conditions were varied confirmed the lognormal form of the mass function and showed that the mean mass of the protostars was similar to the mean initial Jeans mass in the cloud (Klessen \\& Burkert 2000, 2001; Klessen 2001). These calculations enabled us to identify some of the processes that may help to determine the initial mass function.  However, they did not have the resolution to follow the collapsing molecular gas all the way down to the opacity limit for fragmentation (Low \\& Lynden-Bell 1976; Rees 1976; Silk 1977a), or even to resolve the median separation of binary systems of $\\approx 30$ AU \\cite{DuqMay1991}. Thus, we cannot determine the total number of stars that will form or the stellar initial mass function from these calculations, let alone the frequency of binary stars or the importance of star-disc encounters.  This paper presents results from the first calculation to follow the collapse and fragmentation of a large-scale turbulent molecular cloud to form a stellar cluster while resolving beyond the opacity limit for fragmentation. Thus, assuming that fragmentation does not occur at densities greater than those at which the opacity limit sets in (Section \\ref{opacitysec}), it resolves all potential fragmentation, including that which produces binary systems.  This allows us to predict a wide variety of stellar properties.  Two papers that contain results from this calculation have already been published.  They concentrate on the formation mechanisms of brown dwarfs (Bate, Bonnell \\& Bromm 2002a) and close binaries (Bate, Bonnell \\& Bromm 2002b).  In this paper, we consider how the dynamics of star formation determine the properties of stars and brown dwarfs, and we compare these properties with observations.  The outline of this paper is as follows.  In Section 2, we briefly review the opacity limit for fragmentation.  The computational method and the initial conditions for our calculation are described in Section 3.  Section 4 discusses the evolution of the cloud and the star formation that occurs during the calculation. The properties of the resulting stars and brown dwarfs are compared with observations of star-forming regions in Section 5. Finally, in section 6, we give our conclusions ",
        "conclusions": "We have presented the results from one of the most complex hydrodynamical star formation calculations to date.  The calculation follows the collapse and fragmentation of a large-scale turbulent molecular cloud to form a stellar cluster consisting of 50 stars and brown dwarfs.  The opacity limit for fragmentation is mimicked by the use of a non-isothermal equation of state and the resolution is sufficient to resolve all fragmentation before this physical limit is reached.  Binary stars with separations as small as 1 AU and circumstellar discs with radii down to $\\approx 10$ AU are resolved.  The calculation allows us to study the formation mechanisms of stars and brown dwarfs and to determine a wide range of statistical properties for comparison with observation.  We find that star formation is a highly dynamic and chaotic process. A true appreciation of the process can only be obtained from examining an animation of the calculation.  These can be downloaded from http://www.astro.ex.ac.uk/people/mbate or http://www.ukaff.ac.uk/starcluster. Fragmentation occurs both in dense molecular cloud cores and in massive circumstellar discs.  Star-disc encounters form binaries and truncate discs.  Stellar encounters disrupt bound multiple systems.  The star formation occurs on the dynamical timescale of the cloud and the star formation efficiency across the cloud is variable with a low global efficiency ($\\approx 12$\\% when we stop the calculation) but local efficiencies as high as $\\approx 50$\\%.  The high local efficiencies in dense molecular cores result in bursts of star formation because the rapid conversion of gas into stars depletes the high-density gas to such an extent that the star formation essentially comes to a halt until more gas has fallen into the core. When enough new gas has accumulated, another burst of star formation occurs.  We find that the opacity limit for fragmentation sets an initial mass for all fragments of $\\approx 0.005$ M$_\\odot$.  Subsequently, the fragments accrete from the surrounding gas. Those that manage to accrete enough mass become stars ($M\\gsim 0.075$ M$_\\odot$), while the rest are left as brown dwarfs.  We propose that the initial mass function results from this accretion process, with each fragment accreting according to the conditions in which it is formed. The calculation produces a mass function that is consistent with a Salpeter slope ($\\Gamma=-1.35$) above 0.5 M$_\\odot$, a roughly flat distribution ($\\Gamma=0$) in the range $0.006-0.5$ M$_\\odot$, and a sharp cutoff below $\\approx 0.005$ M$_\\odot$.  This is consistent with observational surveys.  Those objects that end up as brown dwarfs stop accreting before they reach stellar masses because they are ejected from the dense gas soon after their formation by dynamical interactions in unstable multiple systems.  Thus, they can be viewed as `failed stars'.  This ejection mechanism is very efficient, producing roughly equal numbers of stars and brown dwarfs (see also Bate et al.\\ 2002a).  However, the close interactions that occur during these dynamical ejections results in a low frequencies ($\\sim 5$\\%) of binary brown dwarf systems. Similarly, the fraction of brown dwarfs with large (radii $\\gsim 20$ AU) circumstellar discs is $\\sim 5$\\%. The accuracy of these frequencies is limited by our small number statistics (for example, we can only exclude a binary brown dwarf frequency of 20\\% at the 94\\% confidence level).  However, further simulations will increase the significance of the predictions. Therefore, observational surveys to determine accurately the frequencies of binary brown dwarfs and the sizes of discs around brown dwarfs should be performed now so that we can test the models.  The calculation produces several binary and higher-order multiple systems.  The opacity limit for fragmentation results in an initial minimum binary separation of $\\approx 10$ AU.  Despite this, we find that 7 close binary systems (separations $< 10$ AU) exist when the calculation is stopped.  These systems are produced by the hardening of initially wider multiple systems through a combination of dynamical encounters, gas accretion, and/or the interaction with circumbinary and circumtriple discs (see also Bate et al.\\ 2002b).  These mechanisms lead to close binaries having a bias towards equal-mass systems and a higher frequency of close binaries for higher-mass stars.  Many of the close binaries also have wider companions.  The resulting frequency of close binary systems is $\\approx 16$\\%, consistent with observations.  Thus, close binary systems need not be formed by fragmentation in situ.  Perhaps the most surprising result of this calculation is that most of the circumstellar discs in the calculation are severely truncated by dynamical encounters.  Most young brown dwarfs should have discs with radii of $\\approx 10$ AU, with large discs ($\\gsim 20$ AU) occurring around only $\\sim 5$\\% of brown dwarfs. The discs around many stars are also severely truncated with the majority having radii $\\lsim 20$ AU (i.e.\\ too small to form our solar system).  Such severe disc truncation, and the associated low masses, may explain the observation that approximately 1/3 of the young stars in Taurus have lost their discs when they are only $\\approx 1$ Myr old (Armitage et al.\\ 2002). Currently, the only star-forming region in which we have information on the size distribution of circumstellar discs is the Orion Trapezium Cluster thanks to the silhouette discs. Our results are consistent with the sizes of discs in the Trapezium Cluster.  However, massive stars are known to be evaporating discs in the Trapezium Cluster and the cluster is much larger than the system we are able to model.  Thus, we strongly encourage observations to determine the size distribution of discs in low-mass star-forming regions such as $\\rho$ Ophiuchus.  This is the first of a new generation of star-formation calculations that resolves all fragmentation and allows us to compare a wide range of statistical properties of stars and brown dwarfs with observations. Future calculations will determine the dependence of these properties on the initial Jeans mass in the cloud, the properties of the turbulence, and will improve the statistical significance of the results. In this way, we hope to understand better the origin of stars and brown dwarfs."
    },
    "0212/astro-ph0212105_arXiv.txt": {
        "abstract": "We use the time-dependent photoionization and dust destruction code developed by Perna \\& Lazzati (2002) to study the time evolution of the medium in a dusty gaseous cloud illuminated by a bright central source that sets on at time zero. We study the case of a bright source, which lasts for a time scale much smaller than the recombination and dust creation time scales. For this reason an equilibrium is never reached. We show that the presence of dust and its properties, such as its composition, can have a big effect on the time scale for the evaporation of the soft X-ray absorption, in particular for ionizing sources with hard spectra. We discuss the profile of evaporation of the soft X-ray absorbing column, as well as how the apparent ionization state of the cloud evolves in time. We finally consider the apparent metallicity of the cloud that is left behind as a function of the cloud and ionizing source properties. ",
        "introduction": "The measurement of the amount of absorption in the soft X-ray band is highly informative about the amount and physical state (e.g. temperature, ionization parameter) of the material that lies along the line of sight between a source and an observer. Observed spectra are usually analyzed under the assumption that the absorber is in equilibrium, either a thermal or photoionization equilibrium (Morrison \\& McCammon 1983; Done et al. 1992; Zdziarski et al. 1995). This assumption is justified if the source that is studied is constant or if its variability time scale is much longer than the electron recombination time scale of the absorbing material. There are however many variable sources in the universe that do not fulfill this condition. Among them it is worth mentioning X-ray transients (see e.g. Schwarz 1973), variable AGNs (see, e.g. Risaliti, Elvis \\& Nicastro 2002) and especially gamma-ray bursts (Perna \\& Loeb 1998; Lazzati \\& Perna 2002, hereafter LP02). In all these cases, out to a certain distance, photoionization acts as a destructive mechanism: each time a photon is absorbed an electron is stripped, and therefore the state of the absorber and its absorption properties are modified.  LP02 (see also Schwarz 1973) analyzed by means of numerical simulations the absorption properties of a gaseous medium suddenly illuminated by a strong photon source with a power-law spectrum. They followed the ionization state of the 12 astrophysically relevant elements and hydrogen producing time resolved transmittance spectra to be compared with the data. In many cases, however, assuming that all the interstellar medium (ISM) is in a gaseous phase is not accurate. There can be molecules (that do not have sizable affect on the X-ray opacity of the medium, see Perna, Lazzati \\& Fiore 2002, hereafter Paper II) and, most importantly, there can be dust grains. These in particular are relevant to X-ray absorption since a large fraction of heavy metals can be condensed in silicate dust grains (Savage \\& Mathis 1979; Laor \\& Draine 1993). The presence of dust grains in the ISM modifies the X-ray absorption properties in two main ways. First, the presence of dust increases the opacity for UV photons, that cannot therefore contribute to the ionization of metals by stripping their external electrons. Most importantly, the metals contained into a dust grain are shielded to photoionization, so that they can survive in an almost neutral state for a much longer time than the same atom in gaseous phase.  We (Perna \\& Lazzati 2002, hereafter Paper I) have developed a code that computes the time-dependent absorption properties of a dusty ISM (possibly enriched with $H_2$ molecules) as a function of time under the illumination of a power-law ionizing continuum that evaporates dust grains, dissociates molecules and ionizes atoms. The code and the physical processes considered are fully described in Paper I. The properties of absorption in the optical range are analyzed in a second paper (Paper II). Here, we focus on the properties of the absorber in the X-ray range, and on how the presence of dust influences the evolution of the opacity under different conditions.  In \\S~2, we briefly introduce the code and describe the outputs that will be used here, in \\S~3 we discuss the main effects of the presence of dust while in \\S~4 we study how these effects depend on the spectrum of the central ionizing source. In \\S~5 we study how the residual column density can be quantified with a single parameter and in \\S~6 we consider the shape of the transmitted spectrum and in particular whether a high ionization parameter can be detected at intermediate stages. In \\S~7, we consider the different ionization time scales of the various elements and the possible effects that this can have on the estimate of metallicity derived from the fit of the relative opacities. Finally, in \\S~8, we summarize and discuss our results.  \\begin{figure} \\psfig{file=f1.eps,width=0.48\\textwidth} \\caption{{The opacity as a function of frequency from optical to X-rays in several time shots after the ionizing continuum onset.  The absorbing medium is a spherical cloud with radius $R=10^{18}$~cm and uniform density $n=10^4$~cm$^{-3}$ ($N_H=10^{22}$~cm$^{-2}$) with the source at its centre.  Abundances are set to the solar value and the dust-to-gas mass fraction is 0.01 (the average Galactic value).} \\label{fig:sim1}} \\end{figure} ",
        "conclusions": "We have analyzed with a dedicated code the time-dependent absorption properties of a cloud subject to a strong photon source that sets in at time $t=0$ at its centre. The novelty of this paper consists in the fact that the presence of dust and its evaporation is fully taken into account by the code, allowing us to study if and how the presence of dust grains can affect the evolution of ionization and vice versa. This has great importance for variable sources such as GRBs and Seyfert galaxies, which are sometimes associated with unusual dust-to-gas ratios (Galama \\& Wijers 2001; Stratta et al. 2002), with variations of the column density (Risaliti et al. 2002; LP02) and with unusual extinction curves (Maiolino et al. 2001).  In this work we concentrated on the X-ray part of the spectrum, since the optical and infrared bands were addressed in a companion paper (Paper II). We study the evolution (evaporation) of the absorption in X-rays of a uniform cloud with solar metallicity for different dust contents, initial column density, cloud radius and hardness of the ionizing spectrum. A source of [1~eV--100~keV] luminosity $L=10^{50}$~erg~s$^{-1}$ is turned on at time $t=0$ at the centre of the cloud.  We first concentrated on the effects that the presence of dust can have on the photoionization rate of metals. We find that, for hard spectrum sources, the presence of dust grains can have a big effect on the photoionization of metals (Fig.~\\ref{fig:dnd}). In fact, if the soft photon flux is not large enough to destroy the dust grain by thermal sublimation, the surviving grains will effectively shield the entrained metals from photoionization (Fig~\\ref{fig:sef}), allowing them to survive in their neutral state for a long time (Fig~\\ref{fig:sim1}). In the more extreme cases, it is possible that the ionizing flux completely ionizes all the atoms in the gaseous phase, leaving behind preferentially large graphite grains, that can contribute a sizable amount of absorption both at optical and X-ray wavelengths.  The final conditions of the cloud depend on both the spectrum, luminosity and variability of the ionizing source as well as on the cloud size and column density (Fig~\\ref{fig:sim1}, \\ref{fig:var}, \\ref{fig:sima}).  We then study two issues related to observations and their modelling. It is in fact customary to model the absorption in the soft X-ray regime by means of equilibrium model for the absorbing medium. In the case we discuss, equilibrium is far from being reached, nevertheless it is useful, in order to exploit real observations, to understand if the predicted spectra look similar to those in equilibrium and how the column density of the absorbing material that would be measured does compare to the real one. We discuss the results of a fit with a cold absorber to our time dependent absorbed spectra (Fig~\\ref{fig:nhs}, \\ref{fig:nhsa}), and then show that a condition in which the absorber seems warm is never reached (Fig~\\ref{fig:xi}, \\ref{fig:xia}).  We finally consider the observed column densities of the various elements, which are time dependent, since atoms completely stripped off their electrons do not contribute to the absorption and are therefore ``invisible''. We find that these ratios are strong functions of time and that as a consequence the metallicities inferred from opacity ratios are not indicative of the pristine composition of the cloud (Fig~\\ref{fig:col}, \\ref{fig:cola}).  Present data do not allow for a detailed comparison with our numerical spectra.  The only evidence of spectral evolution in the early phase of GRBs and afterglows comes from two detections of decreasing column density (Connors \\& Hueter 1998; Frontera et al. 2000), one variable absorption edge (Amati et al. 2000) and a puzzling evidence of variable (but not monotonic) column density (in't Zand et al. 2001). All these detections belong to the $3\\sigma$ realm and do not give much more information than their mere existence. In the near future, thanks to the launch of the Swift satellite in 2003, higher quality spectra will be available, allowing for a more secure detection of transient features and a more meaningful comparison with theoretical spectra."
    },
    "0212/astro-ph0212043_arXiv.txt": {
        "abstract": "We investigate a stationary pair production cascade in the outer magnetosphere of a spinning neutron star. The charge depletion due to global flows of charged particles, causes a large electric field along the magnetic field lines. Migratory electrons and/or positrons are accelerated by this field to radiate gamma-rays via curvature and inverse-Compton processes. Some of such gamma-rays collide with the X-rays to materialize as pairs in the gap. The replenished charges partially screen the electric field, which is self-consistently solved together with the energy distribution of particles and gamma-rays at each point along the field lines. By solving the set of Maxwell and Boltzmann equations, we demonstrate that an external injection of charged particles at nearly Goldreich-Julian rate does not quench the gap but shifts its position and that the particle energy distribution cannot be described by a power-law. The injected particles are accelerated in the gap and escape from it with large Lorentz factors. We show that such escaping particles migrating outside of the gap contribute significantly to the gamma-ray luminosity for young pulsars and that the soft gamma-ray spectrum between 100~MeV and 3~GeV observed for the Vela pulsar can be explained by this component. We also discuss that the luminosity of the gamma-rays emitted by the escaping particles is naturally proportional to the square root of the spin-down luminosity. ",
        "introduction": "\\label{sec:intro} Recent years have seen a renewal of interest in the theory of particle acceleration in pulsar magnetospheres, after the launch of the {\\it Compton Gamma-ray Observatory} (e.g., for the Vela pulsar, Kanbach et al. 1994, Fierro et al. 1998; for PSR B1706--44, Thompson et al. 1996; for Geminga, Mayer-Hasselwander et al. 1994, Fierro et al. 1998; for PSR B1055--52, Thompson et al. 1999). The modulation of the $\\gamma$-ray light curves testifies to the particle acceleration either at the polar cap (Harding, Tademaru, \\& Esposito 1978; Daugherty \\& Harding 1982, 1996; Sturner, Dermer, \\& Michel 1995), or at the vacuum gaps in the outer magnetosphere (Cheng, Ho, \\& Ruderman 1986a,b, hereafter CHR; Chiang \\& Romani 1992, 1994; Romani \\& Yadigaroglu 1995; Higgins \\& Henriksen 1997, 1998). Both models predict that electrons and positrons are accelerated in a charge depletion region, a potential gap, by the electric field along the magnetic field lines to radiate high-energy $\\gamma$-rays via the curvature process. However, there is an important difference between these two models: An polar-gap accelerator releases very little angular momenta, while an outer-gap one could radiate them efficiently. In addition, three-dimensional outer-gap models commonly explain double-peak light curves with strong bridges observed for $\\gamma$-ray pulsars. On these grounds, the purpose of the present paper is to explore further into the analysis of the outer-gap accelerator.  In the CHR picture, the gap is assumed to be geometrically thin in the transfield direction on the poloidal plane in the sense $D_\\perp \\ll W$, where $D_\\perp$ represents the typical transfield thickness of the gap, while $W$ does the width along the magnetic field lines. In this limit, the acceleration electric field is partially screened by the zero-potential walls separated with a small distance $D_\\perp$; as a result, the gap, which is assumed to be vacuum, extends from the null surface to (the vicinity of) the light cylinder. Here, the null surface is defined as the place on which the local Goldreich-Julian charge density \\begin{equation} \\rhoGJ \\equiv -\\frac{\\Omega B_z(s,z)}{2\\pi c} \\label{eq:def_rhoGJ} \\end{equation} vanishes, where $\\Omega$ refers to the angular frequency of the neutron star, $B_z$ the magnetic field component projected along the rotational axis, and $c$ the velocity of light; $s$ and $z$ refer to the coordinates parallel and perpendicular, respectively, to the poloidal magnetic field. The star surface corresponds to $s=0$; $s$ increases outwardly along the field lines. The last-open field line corresponds to $z=0$; $z$ increases towards the magnetic axis (in the same hemisphere).  If $B_z>0$ holds in the starward side of the null surface, a positive acceleration field arises in the gap. The light cylinder is defined as the surface where the azimuthal velocity of a plasma would coincide with $c$ if it corotated with the magnetosphere. Its radius from the rotational axis becomes the so-called \\lq light cylinder radius', $\\rlc \\equiv c / \\Omega$. Particles are not allowed to migrate inwards beyond this surface because of the causality in special relativity.  It should be noted that the null surface is not a special place for the gap electrodynamics in the sense that the plasmas are not completely charge-separated in general and that the particles freely pass through this surface inwards and outwards. Therefore, the gap inner boundary is located near to the null surface, not because particle injection is impossible across this surface (as previously discussed), but because the gap is vacuum and transversely thin.  Then what occurs in the CHR picture if the gap becomes no longer vacuum? To consider this problem rigorously, we have to examine the Poisson equation for the electrostatic potential. In fact, as will be explicitly demonstrated in the next section, the original vacuum solution obtained in the pioneering work by CHR cannot be applied to a non-vacuum CHR picture. We are, therefore,  motivated by the need to solve self-consistently the Poisson equation together with the Boltzmann equations for particles and $\\gamma$-rays. Although the ultimate goal is to solve three-dimensional issues, a good place to start is to examine one-dimensional problems. In this context, Hirotani and Shibata (1999a,~b,~c; hereafter Papers~I,~II,~III) first solved the Boltzmann equations together with the Maxwell equations one-dimensionally along the field lines, extending the idea originally developed for black-hole magnetospheres by Beskin et al. (1992). In Paper~I, II, and III, they assumed that the gap is geometrically thick in the transfield direction in the sense $D_\\perp \\gg W$.  There is one important finding in this second picture: The gap position shifts if there is a particle injection across either of the boundaries (Hirotani \\& Shibata 2001, 2002a,b; hereafter Papers~VII, VIII, IX). For example, when the injection rate across the outer (or inner) boundary becomes comparable to the typical Goldreich-Julian value, the gap is located close to the neutron star surface (or to the light cylinder). In other words, an outer gap is not quenched even when the injection rate of a completely charge-separated plasma across the boundaries approaches the typical Goldreich-Julian value. Thus, an outer gap can coexist with a polar-cap accelerator; this forms a striking contrast to the first, CHR picture. It is also found in the second picture that an outer gap is quenched if the {\\it created} particle density within the gap exceeds several percent of the Goldreich-Julian value. That is, the {\\it discharge} of created pairs is essential to screen the acceleration field.  The purpose of this paper is to examine the second picture more closely. In all the previous works in the second picture, the particle energy distribution has been assumed to be mono-energetic in the sense that the particles attain the equilibrium Lorentz factor at each point, in a balance between the electrostatic acceleration and the radiation-reaction forces. In the present paper, we discard this assumption and explicitly consider the energy dependence of particles by solving the Boltzmann equations for positrons and electrons. We will demonstrate that the particle energy distribution cannot be represented either by a power law or by the mono-energetic approximation. We will further show that a soft power-law spectrum is generally formed in 100~MeV-GeV energies as a result of the superposition of the curvature spectra emitted by particles migrating at different positions.  In the next section, we describe the difficulties of electrodynamics found in the first picture. We then present the basic equations in \\S~\\ref{sec:basic}. and apply the theory to four $\\gamma$-ray pulsars and compare the predictions with observations in \\S~\\ref{sec:app}. In the final section, we discuss the possibilities of the unification of our picture with the CHR picture, as well as the unification of the outer-gap and polar-cap models. ",
        "conclusions": "\\label{sec:discussion} In summary, we have quantitatively examined the stationary pair-production cascade in an outer magnetosphere, by solving the set of Maxwell and Boltzmann equations one-dimensionally along the magnetic field. We revealed that an accelerator (or a potential gap) is quenched by the created pairs in the gap but is {\\it not} quenched by the injected particles from outside of the gap, and that the gap position shifts as a function of the injected particle fluxes: If the injection rate across the inner (or outer) boundary approaches the typical Goldreich-Julian value, the gap is located near to the light cylinder (or the star surface). It should be emphasized that the particle energy distribution is not represented by a power law, as assumed in some of previous outer-gap models. The particles escape from the gap with sufficient Lorentz factors and emit significant photons in 100~MeV--3~GeV energies via curvature radiation outside of the gap. The $\\gamma$-ray spectrum including this component explains the phase-averaged EGRET spectra for the Vela pulsar, the Geminga pulsar and PSR~B1055-52 between 100~MeV and 6~GeV. TeV fluxes are unobservable with current ground-based telescopes for these pulsars.  We show that synchro-curvature process can be approximated by the pure-curvature one in the next subsection. We then point out an implication to the $\\gamma$-ray luminosity versus the spin-down luminosity in \\S~\\ref{sec:lumin}, and discuss future extensions of the present method in \\S\\S~\\ref{sec:return}--\\ref{sec:unif_pol}.  \\subsection{Synchrotron vs. Curvature Processes} \\label{sec:small_pitch} Let us first examine the pitch angles of particles injected across the inner boundary. Imposing that the synchrotron cooling length scale, \\begin{equation} l_{\\rm sync} \\equiv \\frac{\\Gamma m_{\\rm e} c^2}{P_{\\rm sync}/c} = \\frac{3 m_{\\rm e}^3 c^6}{2 e^4 B^2 \\Gamma \\sin^2\\chi} \\label{eq:cool_sync} \\end{equation} is greater than the distance, $\\varepsilon \\rlc$, for the particles to migrate before arriving the outer gap, we obtain \\begin{equation} \\sin\\chi < 7.2 \\times 10^{-6} \\left( \\frac{\\Omega_2}{\\varepsilon_{0.1}\\Gamma_2} \\right)^{1/2} B_{10}^{-1}, \\label{eq:cool_pitch} \\end{equation} where $\\varepsilon_{0.1}\\equiv \\varepsilon/0.1$, $\\Gamma_2 \\equiv \\Gamma/10^2$, $\\Omega_2 \\equiv \\Omega/(10^2 \\mbox{rad s}^{-1})$, and $B_{10} \\equiv B/(10^{10} \\mbox{G})$. Particles having initial pitch angles greater than this value will lose transverse momentum to reduce the pitch angles below this value, while migrating the distance $\\varepsilon \\rlc$. Therefore, if the particles are supplied from the polar cap with $\\Gamma \\sim 10^2$ for instance, we can safely state that the the pitch angles are less than $10^{-5}$ rad.  If a particle having such a small pitch angle is accelerated to $\\Gamma=10^7$ in a weak magnetic field region $B=10^6$~G, the ratio between the synchrotron and the curvature radiation rate becomes \\begin{eqnarray} \\lefteqn{ \\frac{P_{\\rm sync}}{\\Pcv} = \\left( \\frac{\\rho_{\\rm c} \\sin\\chi} {\\Gamma m_{\\rm e}c^2/eB} \\right)^2} \\nonumber\\\\ &=& 7.7 \\times 10^{-3} \\Omega_2^{-2} B_6^2 \\Gamma_7^{-2} \\left(\\frac{\\rho_{\\rm c}}{0.5\\rlc} \\cdot \\frac{\\sin\\chi}{10^{-5}} \\right)^2, \\label{eq:PsyPcv} \\end{eqnarray} where $B_6 \\equiv B/10^6 \\mbox{G}$, $\\Gamma_7 \\equiv \\Gamma/10^7$. It follows that the synchro-curvature radiation can be approximated by a pure-curvature one for the particles injected across the inner boundary. It is noteworthy that the density of created particles is much less than that of injected ones (i.e., $j_{\\rm gap} \\ll j^{\\rm in}$); therefore, the synchro-curvature effects for the created particles, which have much greater pitch angles compared with the injected ones, are negligibly small for the gap electrodynamics as well for the resultant $\\gamma$-ray spectrum.  \\subsection{Gamma-ray vs. Spin-down Luminosities} \\label{sec:lumin} It should be noted that the emission from the escaping particles attain typically $40\\%$ of the total $\\gamma$-ray luminosity for young pulsars. Thus, it is worth mentioning its relationship with the spin-down luminosity, \\begin{equation} L_{\\rm spin}= -I \\Omega \\dot{\\Omega}  \\propto \\Omega^{n+1}, \\label{eq:spindown} \\end{equation} where the braking index $n$ is related to the spin-down rate as \\begin{equation} \\dot{\\Omega}= -k \\Omega^n. \\label{eq:braking} \\end{equation} If the spin down is due to the magnetic dipole radiation, we obtain $n=3$.  The outwardly propagating particles escape from the gap with spatial number density (eq.~[\\ref{eq:def-n}]) \\begin{equation} N_{\\rm out} = (j^{\\rm in}+j_{\\rm gap})\\frac{\\Omega B^{\\rm out}}{2\\pi ce}, \\label{eq:NeOUT} \\end{equation} where $B^{\\rm out}=B(\\xi^{\\rm out})$. Therefore, the energy carried by the escaping particles per unit time is given by \\begin{equation} L_{\\rm esc} = D_\\perp^2 c N_{\\rm out} \\Gamma_{\\rm esc} m_{\\rm e}c^2, \\label{eq:def_Lesc} \\end{equation} where $\\Gamma_{\\rm esc} (\\sim 10^{7.5})$ refers to the Lorentz factor of escaping particles. Note that $\\Gamma_{\\rm esc}$ is essentially determined by the equilibrium Lorentz factor (dotted line in fig.~\\ref{fig:EVela_75a}) near the gap center. Since the equilibrium Lorentz factor depends on the one-fourth power of $\\Ell$, the variation of $\\Gamma_{\\rm esc}$ on pulsar parameters is small. We can approximate $B^{\\rm out}$ as \\begin{equation} B^{\\rm out} \\sim \\frac{\\mu}{\\rlc^3} \\left( \\frac{\\rlc}{r^{\\rm out}} \\right)^3, \\label{eq:Bout} \\end{equation} where $r^{\\rm out}$ refers to the distance of the outer boundary of the gap from the star center. Let us assume that the position of the gap with respect to the light cylinder radius, $r^{\\rm out}/\\rlc$, does not change as the pulsar evolves; this situation can be realized if $j^{\\rm in}-j^{\\rm out}$ is unchanged. Evaluating $B$ at $r=0.5\\rlc$, we obtain \\begin{eqnarray} L_{\\rm esc} &=& \\frac{4 \\Gamma_{\\rm esc} m_{\\rm e}c}{\\pi e} \\mu \\Omega^2 \\left(\\frac{D_\\perp}{\\rlc}\\right)^2 \\nonumber\\\\ &\\propto& L_{\\rm spin}{}^{0.5}, \\label{eq:Lesc} \\end{eqnarray} where $n=3$ is assumed in the second line. To derive this conclusion, it is essential that the particles are not saturated at the equilibrium Lorentz factor. Thus, the same discussion can be applied irrespective of the gap position or the detailed physical processes involved. For example, an analogous conclusion was derived for a polar-cap model by Harding, Muslimov, and Zhang (2002). It is, therefore, concluded that the observed relationship $L_\\gamma \\propto L_{\\rm spin}{}^{0.5}$ merely reflects the fact that the particles are unsaturated in the gap and does not discriminate the gap position.  Let us compare this result with what would be expected in the CHR picture. Since the gap is extended significantly along the field lines in the CHR picture, particles are saturated at the equilibrium Lorentz factor to lose most of their energies within the gap, rather than after escaping from it. We can therefore estimate the $\\gamma$-ray luminosity as \\begin{equation} L_{\\rm gap} = (D_\\perp D_\\phi W) \\cdot N_{\\rm out} \\cdot \\Pcv \\label{eq:def_Lgap} \\end{equation} where $D_\\phi$ refers to the azimuthal thickness of the gap, $\\Pcv$ [ergs s$^{-1}$ (particle)$^{-1}$] represents the curvature radiation rate. Noting that the particle motion saturates at the equilibrium Lorentz factor satisfying $\\Pcv/c=e(-d\\Psi/ds)$, recalling that the acceleration field is given by $-d\\Psi/ds \\approx \\Omega B D_\\perp^2 / 4\\rho_{\\rm c}c$ in the CHR picture, and evaluating $B$ at $r=\\rlc$, we obtain \\begin{eqnarray} L_{\\rm gap} &=& \\frac{\\mu^2 \\Omega^4}{4\\pi c^3} \\frac{D_\\perp^3 D_\\phi W}{\\rlc^5} \\left(\\frac{\\rho_{\\rm c}}{0.5\\rlc}\\right)^{-1} \\nonumber\\\\ &\\propto& L_{\\rm spin} \\label{eq:Lgap}, \\end{eqnarray} where $n=3$ is assumed again in the second line. Even though the escaping particles little contribute to the $\\gamma$-ray luminosity in the CHR picture, it is worth mentioning the work done by Crusius--W$\\ddot{\\rm a}$tzel and Lesch (2002), who accurately pointed out the importance of the escaping particles in the CHR picture, when we interpret $L_\\gamma \\propto L_{\\rm spin}^{0.5}$ relation.  As we have seen, the particles being no longer accelerated contribute for the $\\gamma$-ray luminosity that is proportional to $L_{\\rm spin}^{0.5}$. Reminding that the particles migrate with larger Lorentz factors than the equilibrium value in the outer part of the gap (see fig.~\\ref{fig:EVela_75a}), we can expect roughly half of the $\\gamma$-ray luminosity is proportional to $L_{\\rm spin}^{0.5}$ (mainly between 100~MeV and 1~GeV), and the rest of the half to $L_{\\rm spin}$ (mainly above 1~GeV). As a pulsar ages, its declined surface emission results in a large pair-production mean free path, and hence $W$. Because $\\vert \\rhoGJ \\vert \\propto r^{-3}$ becomes small in the outer part of such an extended gap, $\\Ell(\\xi)$ deviates from quadratic distribution to decline gradually in the outer part (fig.~\\ref{fig:EGemi_60}). As a result, particles tend to be saturated at the equilibrium value. On these grounds, we can predict that the $\\gamma$-ray luminosity tends to be proportional to $L_{\\rm spin}$ with age, deviating from $L_{\\rm spin}^{0.5}$ dependence for young pulsars.  In the present paper, we have examined the set of Maxwell and Boltzmann equations one-dimensionally both in the configuration and the momentum spaces (i.e., only $\\xi$ and $\\Gamma$ dependences are considered.) In the next three sections, we discuss the extension of the present method into higher dimensions  \\subsection{Returning Particles} \\label{sec:return} If we consider the pitch-angle dependence of particle distribution functions, we can compute the radiation spectrum with synchro-curvature formula (Cheng and Zhang 1996). Moreover, we can also consider the returning motion of particles inside and outside of the gap. The returning motion becomes particularly important when both signs of charge are injected across the boundary. For example, not only positrons but also electrons could be injected across the inner boundary from the polar-cap accelerator. If $\\Ell>0$ for instance, the injected electrons return in the gap. This returning motion significantly affects the Poisson equation, if their injection rate is a good fraction of the Goldreich-Julian value.  It remains an unsettled issue whether an outer-gap accelerator resides on the field lines on which a polar-cap accelerator exists. To begin with, let us consider the case when the plasma flowing between the polar cap and the outer-gap accelerator is completely charge separated. Such a situation can be realized, for instance, if only positively charged particles are ejected outwardly from the polar cap while there is virtually no electrons ejected inwardly from the outer gap. Neglecting the pair production, current conservation law gives the charge density, $\\rho_{\\rm e}$, per unit magnetic flux tube as \\begin{equation} \\frac{\\rho_{\\rm e}}{B} \\propto \\frac{j_{\\rm tot}}{v}, \\label{eq:sep_flow} \\end{equation} where $v$ refers to the particle velocity along the field line, and $j_{\\rm tot}$ the conserved current density per magnetic flux tube. At each point along the field line, $\\rho_{\\rm e}$ should match $\\rho_{\\rm GJ}$. If the field line intersects the null surface, $\\rho_{\\rm e}$ must vanish there; this obviously violates the causality in special relativity. Therefore, a stationary ejection of a completely charge-separated plasma from the polar cap can be realized only along the field lines between the magnetic axis and those intersecting the null surface at the light cylinder. On these grounds, it was argued that an outer-gap accelerator, which is formed close to the last-open field line, may not resides on the same field lines on which a polar-cap accelerator resides. This has been, in fact, the basic idea that an outer gap will not be quenched, because the particles ejected from the polar cap will flow along the different field lines. This idea was welcomed in outer-gap models, because a gap has been considered to be quenched if the external particle injection rate becomes comparable to the Goldreich-Julian value, which was proved to be incorrect in this paper.  In general, however, the plasmas are not completely charge separated and consist of both signs of charge (e.g., positrons and electrons). Such a situation can be realized, for instance, if both charges are ejected outwardly from a polar-cap accelerator, or if positively charged particles are ejected outwardly from the polar cap while electrons are ejected inwardly from the outer gap, or if there is a pair production between the two accelerators. In these cases, the velocities of both charges will be adjusted so that both the current conservation and $\\rho_{\\rm e}=\\rho_{\\rm GJ}$ are satisfied at each point along the field lines. Therefore, it seems likely that a polar-cap accelerator and an outer-gap accelerator reside on the same field lines.  To examine if there is a stationary plasma flow between the polar cap and the outer gap, we must extend the present analysis into two dimensional momentum space in the sense that the pitch-angle dependence of the particle distribution functions is taken into account in addition to the Lorentz factor dependence. For example, if both charges are ejected from the polar-cap accelerator, electrons will return in the outer gap, screening the original acceleration field in the gap, and violating the original balance of $\\rho_{\\rm e}=\\rhoGJ$ outside of the gap. Because the returning motion of particles can be treated correctly if we consider the pitch-angle evolution of the distribution functions, and because the pair production is already taken into account, our present method is ideally suited to investigate the plasma flows and $\\Ell$ distribution self-consistently inside and outside of the gap.  \\subsection{Unification of Outer-gap Models} \\label{sec:unif_out} In addition to the extension into a higher dimensional momentum space, it is also important to extend the present method into a two- or three-dimensional configuration space. In particular, determination of the perpendicular thickness, $D_\\perp$, is important to constrain gap activities. There have been, in fact, some attempts to constrain $D_\\perp$ in the CHR picture. Since $-d\\Psi/ds$ is proportional to $B D_\\perp^2$, particles energies, and hence the $\\gamma$-ray energies increase with increasing $D_\\perp$ (for a fixed $B$). Zhang and Cheng (1997) constrained $D_\\perp$, by considering the condition that the $\\gamma$-rays cause photon-photon pair production in the gap. Subsequently, Cheng, Ruderman, and Zhang (2000) extended this idea into three-dimensional magnetosphere and discussed phase-resolved $\\gamma$-ray spectra for the Crab pulsar. In addition, Romani (1996) discussed the evolution of the $\\gamma$-ray emission efficiency and computed the phase-resolved spectra for the Vela pulsar, by assuming that $BD_\\perp^2$ declines as $r^{-1}$. However, in these works, screening effects due to pair production has not been considered; thus, the obtained $D_\\perp$, as well as the assumed gap position along the magnetic field, are still uncertain.  On the other hand, in our approach (picture), $D_\\perp$ is not solved but only adjusted so that the $\\gamma$-ray flux may match the observations. Therefore, the question we must consider next is to solve such geometrical and electrodynamical discrepancies between these two pictures. We can investigate this issue by extending the present method into higher spatial dimensions.  \\subsection{Unification of Outer-gap and Polar-cap Models} \\label{sec:unif_pol} Electrodynamically speaking, the essential difference between outer-gap and polar-cap accelerators is the value of the optical depth for pair production. In an outer-gap accelerator, pair production takes place via $\\gamma$-$\\gamma$ collisions and its mean-free path is much greater than the light cylinder radius. Therefore, a pair production cascade takes place gradually in the gap. In such a gap, $\\Ell$ is screened out by the \\lq generalized Goldreich-Julian charge density', $f_{\\rm null}$ (eq.~[\\ref{eq:def_fnull}]), which increases outwards if $\\mbox{\\boldmath$\\Omega$}\\cdot\\mbox{\\boldmath$B$}>0$. Since $f_{\\rm null}$ is negative (or positive) in the inner (or outer) part of the gap, there is a surface on which the right-hand side of equation~(\\ref{eq:Poisson_1Db}) vanishes, as long as the injected current density is less than the Goldreich-Julian value. The gap is located around this \\lq generalized null surface', which is explicitly defined in \\S~2 of Paper~VII.  On the contrary, in a polar-cap accelerator, pair production takes place mainly via magnetic pair production, of which mean free path is much less than the star radius for a typical magnetic field strength ($B \\sim 10^{12}$~G, say). As a result, a pair production avalanche takes place in a limited region, which is called as the \\lq pair formation front', in the gap (Fawley, Arons, \\& Sharlemann 1977; Harding \\& Muslimov 1998, 2001, 2002; Shibata, Miyazaki, Takahara 1998, 2002; Harding, Muslimov, Zhang 2002). In the pair formation front, a small portion of the particles return to screen out $\\Ell$. Such a returning motion can be self-consistently solved together with $\\Ell$ by our present method, if we implement the magnetic pair production and the resonant IC scattering redistribution functions in the source terms of the particles' and $\\gamma$-rays' Boltzmann equations. We can execute the same advection-phase computation in CIP scheme; thus, all we have to do is to add these source terms in the non-advection-phase computation, which is not very difficult. Since analogous boundary conditions (e.g., $\\Ell=0$ for a space-charge limited flow) will be applied, we expect the present method is also applicable to a polar-cap accelerator. This is an issue to be examined in our subsequent papers.   \\par \\vspace{1pc}\\par  One of the authors (K. H.) wishes to express his gratitude to Drs. K. Shibata and A. Figueroa-Vin{$\\bar{\\rm a}$}s for valuable advice on numerical analysis, and to Drs. K.~S.~Cheng and C. Thompson for fruitful discussion on theoretical aspects. He also thanks Canadian Institute for Theoretical Astrophysics for welcoming him as a visiting researcher."
    },
    "0212/astro-ph0212261_arXiv.txt": {
        "abstract": "This study calls attention to the importance of properly coupling the molecular opacities to the actual surface abundances of TP-AGB stars that experience the third dredge-up and/or hot-bottom burning, i.e. with surface abundances of carbon and oxygen varying with time. New TP-AGB calculations with variable opacities -- replacing the usually adopted solar-scaled opacity tables -- have proven to reproduce, for the first time, basic observables of carbon stars, like their effective temperatures, C/O ratios, and near-infrared colours. Moreover, it turns out that the effect of envelope cooling -- due to the increase in molecular opacities -- may cause other important effects, namely: i) shortening of the C-star phase; ii) possible reduction or shut-down of the third dredge-up in low-mass carbon stars; and iii) weakening or even extinction of hot-bottom burning in intermediate-mass stars. ",
        "introduction": "The observed spectral dichotomy between M-type (with C/O$<1$) and C-type (with C/O$>1$) stars was first explained by Russell (1934) on the basis of of molecular equilibria calculations. It was pointed out that the C/O ratio plays the key-role in determining the different patterns of molecular abundances, hence different dominant molecular bands, characterising the two classes of giant stars.  Despite this basic fact, the  description of molecular opacities is still improper in  most evolution models of AGB stars. In fact, the usually adopted opacity tables (e.g. Alexander \\& Ferguson 1994) are strictly valid for solar-scaled abundances of elements heavier than helium, corresponding to C/O$ = 0.48$ (hereinafter also $\\kappa_{\\rm fix}$ prescription).  Therefore, it is already clear that the inadequacy of the opacity prescription becomes  particularly serious when modelling carbon stars, characterised by surface C/O$ > 1$ as a consequence of recurrent third dredge-up episodes during the TP-AGB evolution.  Marigo (2002) investigates the effects on the AGB evolution due to variable molecular opacities (hereinafter also $\\kappa_{\\rm var}$ prescription), that are now computed consistently to the current envelope chemical composition of TP-AGB models experiencing the third dredge-up and hot-bottom burning (HBB).  As illustrated in the following, from comparing the new $\\kappa_{\\rm var}$ with the standard  $\\kappa_{\\rm fix}$ results, the impact turns out indeed significant. \\begin{figure} \\plotfiddle{marigof1.ps}{2.8in}{0.}{42.}{42.}{-135}{-75} \\caption{Molecular abundances (in terms of partial  pressures) of a few atomic and molecular species as a function of the C/O ratio, assuming  a gas pressure $P_{\\rm gas} = 10^3$ dyne cm$^{-2}$, and a temperature $T = 2500$ K. The vertical line marks the molecular concentrations for a solar composition, with C/O$\\sim 0.48$ } \\label{fig_molcovar} \\end{figure} ",
        "conclusions": "This explorative study has shown how large is the impact of introducing variable molecular opacities in AGB models. This has improved the comparison with observations and brought  many new results, which may critically change and revitalise, in various aspects, the present scenario of AGB evolution."
    },
    "0212/astro-ph0212057_arXiv.txt": {
        "abstract": "{ We present evolutionary models of zero-metallicity very massive objects, with initial masses in the range 120 $M_{\\odot}$ -- 1000 $M_{\\odot}$, covering their quiescent evolution up to central carbon ignition. In the attempt of exploring the possible occurrence of mass loss by stellar winds, calculations are carried out with recently-developed formalisms for the  mass-loss rates driven by radiation pressure (Kudritzki 2002) and stellar rotation (Maeder \\& Meynet 2000). The study completes the previous analysis by Marigo et al. (2001) on the constant-mass evolution of primordial stars. Our results indicate that radiation pressure (assuming a minimum metallicity $Z = 10^{-4}\\times Z_{\\odot}$) is not an efficient driving force of mass loss, except for very massive stars with $M \\ga 750 \\, M_{\\odot}$. On the other hand, stellar rotation might play a crucial role in triggering  powerful stellar winds, once the $\\Omega\\Gamma$-limit is approached. However, this critical condition of intense mass loss can be maintained just for short, as the loss of angular momentum due to mass ejection quickly leads to the spinning down of the star. As by-product to the present work, the wind chemical yields from massive zero-metallicity stars are presented. The helium and metal enrichments, and the resulting $\\Delta Y/\\Delta Z$ ratio are briefly discussed. ",
        "introduction": "\\label{sec_intro}  The first generation of stars ever born in the Universe (Population III) is assigned a role of paramount importance for several astrophysical issues. In particular, the formation of a primeval population of very massive objects (VMOs, with initial masses in the range $10^2-10^5 \\, M_{\\odot}$) was invoked in the '80s to explain several questions, such as the observed metallicities of Population II stars, the primordial helium abundance, the re-ionisation of cosmic matter after the Big-Bang, the missing mass in clusters of galaxies and galactic haloes, the G-dwarf problem (see Carr et al. 1984 for an extensive review). In those years a few evolutionary calculations of zero-metallicity VMOs were carried out (e.g. Bond et al. 1984, El Eid et al. 1983, Ober et al. 1983, Klapp 1984ab), but afterwards the scientific production on the evolution of the first stars dropped off.  The recently renewed interest in Pop-III stars, essentially driven by the nowadays flourishing of observations at very low metallicity and/or high redshift (refer to e.g. the proceedings of the ESO symposium edit by Weiss et al. (2000); see also Kudritzki et al. (2000), Bromm et al. (2001), Schaerer (2002), Panagia (2002) for extensive analyses on the expected observable properties of primordial stellar populations) has again stimulated  the calculation of stellar structures made evolve with initial metal-free chemical composition (e.g. Cassisi et al. 2001; Marigo et al. 2001 and references therein).  Marigo et al. (2001; thereinafter Paper I) presented an extensive study of the evolutionary properties of zero-metallicity stars over a wide range of initial masses  ($0.7 \\, M_{\\odot} \\la M \\la 100 M_{\\odot}$). In that work -- to which the reader is referred for all the details -- we calculated all stellar tracks at constant mass.  Our study on Pop-III stars is now extended to very massive objects (VMOs), with initial masses in the range $120-1000\\, M_{\\odot}$, that is chosen not only for continuity with Paper I, but also in consideration of current theoretical indications that the primordial initial mass function  (IMF) might have peaked in the (very) high-mass domain (with $M \\ga 100 \\, M_{\\odot}$; see e.g. Bromm et al. 1999, 2002; Nakamura \\& Umemura 2001; Abel et al. 2000).  To this aim we calculate evolutionary models for zero-metallicity VMOs that cover the major core-burning phases, extending from hydrogen to carbon ignition. With aid of available analytic formalisms, we address the question of the possible occurrence of mass loss via stellar winds, which is an important, but still problematic, aspect of stellar evolution at zero metallicity.  In massive stars with ``normal'' chemical composition the principal driving force resides in the capability of metallic ions, present in the atmosphere, to absorb radiative momentum and transfer it to the outermost layers, that can be then accelerated beyond their escape velocity (Castor et al. 1975; Pauldrach et al. 1986). In a gas composed only of hydrogen and helium, like the primordial one, the lack of metallic ions sets the first important difference and the question arises: Is the radiative acceleration due to the H and He lines strong enough to trigger significant mass loss ? Furthermore,  may other possible processes -- e.g. related to pulsation instability and stellar rotation -- be efficient mechanisms to drive mass loss from Pop-III massive stars ?  Some of these questions have already been addressed by other investigations, others still deserve to be quantitatively analysed and discussed. There are few studies in the literature that present evolutionary models  of very massive zero-metallicity stars with mass loss  during the pre-supernova phases. In the very massive domain ($500-1000\\, M_{\\odot}$) Klapp (1983, 1984) carried out evolutionary calculations by adopting a simple empirical law for mass-loss (Barlow \\& Cohen 1977) -- based on Galactic observations -- that linearly scales with the stellar luminosity. On the basis of semi-analytical models of very massive objects (with masses $10^2 -10^5\\, M_{\\odot}$), Bond et al. (1984) argued that the same kind of dynamical instability arising in the H-burning shell of Population I models -- which should cause the ejection of the entire  envelope --  might also affect Population III VMOs, depending on the actual abundance of the CNO catalysts in the H-shell. The evolution of zero-metallicity stars (with masses $80 - 500\\, M_{\\odot}$) during  the nuclear phases of H- and He-burning was calculated by El Eid et al. (1983) with the adoption of a semi-empirical relation to derive  the mass loss rate (Chiosi 1981). Recently Baraffe et al. (2001, see also Heger et al. 2001) performed a linear stability analysis on metal-free very massive models (with masses $120-500\\, M_{\\odot}$), and  investigated the related possibility that mass loss may characterise the pulsation-unstable stages. Heger \\& Woosley (2002) calculated the evolution and nucleosynthesis of quite massive Population III stars (with masses $\\sim 100 - 300\\, M_{\\odot}$) including the final fate of the supernova event, but no mass-loss is assumed to occur during the hydrostatic phases of major nuclear burnings.  In the present work we attempt to further explore the issue of mass loss in primordial conditions, by carrying out new evolutionary calculations of  zero-metallicity VMO  with recently developed mass-loss formalisms, related to radiation pressure and stellar rotation.  The paper is organised as follows. Section \\ref{sec_evolcalc} introduces new evolutionary calculations for zero-metallicity VMOs, having initial masses of 120, 250, 500, 750, and 1000 $M_{\\odot}$. Mass loss is included with  the adoption  of recent analytic formalisms, namely: Kudritzki (2002) for radiation-driven mass loss, and Maeder \\& Meynet (2000) for the additional effect of stellar rotation. General characteristics of the models are analysed in relation to energetics, nuclear lifetimes, internal structure, surface and chemical properties. The efficiency of mass loss is discussed on the basis of the adopted formalisms. The corresponding predictions for the chemical yields, ejected via stellar winds, are given in Sect.~\\ref{sec_yields}. In particular, the expected helium and metal enrichment, produced by a hypothetical burst of Pop-III star formation, is derived and discussed as a function of the involved parameters. In Sect.~\\ref{sec_concl} we express a few concluding remarks. Finally, stellar isochrones for very young ages are presented in Sect.~\\ref{sec_tableisoc}. ",
        "conclusions": "\\label{sec_concl}  In this work we have discussed the evolutionary properties of zero-metallicity very massive objects ($120\\, M_{\\odot} \\le M \\le 1000\\, M_{\\odot}$), on the basis of new calculations that extend the work presented in Paper I. Stellar isochrones (see Appendix A) for ages from 16~Gyr down to $10^{4}$~yr are available at the web-address http://pleiadi.pd astro.it.  In the attempt to estimate the possible effects produced by stellar winds from primordial VMOs, we adopt recent formalisms that describe the role of radiation pressure and stellar rotation as mass-loss driving factors. The  emerging picture is the following:  \\begin{itemize} \\item At extremely low metallicity ($Z \\approx 10^{-6}$), the mechanism of line-radiation transfer should be scarcely efficient, except for very large masses, say $\\ga 750 M_{\\odot}$. We also find that, as long as the mass loss front penetrates into the chemical profile left by the convective core during the H-burning phase, the maximum CNO abundance exposed at the surface does not exceed $\\approx 10^{-9}-10^{-8}$, which are typical values required by the CNO-cycle operating in stars with original metal-free composition. Such degree of chemical self-pollution is also too low to trigger efficient radiation-driven winds from zero-metallicity massive stars. Instead, much larger surface enrichment may be attained whenever the nuclear products of He-burning, like carbon and oxygen -- potentially able to produce large mass-loss rates -- were  brought up to the surface by either the penetrating mass-loss front, or some dredge-up process. In our computations this occurs only for the $1000\\, M_{\\odot}$ model.  \\item Our calculations also indicate that rotating very massive  stars may actually reach  the critical condition defined as the $\\Omega\\Gamma$-limit, which should be likely accompanied by large mass-ejection rates (here taken as large as $\\dot M_{\\rm crit} = 10^{-3}$ yr$^{-1}$). However, the net impact on mass loss should be quite limited, given the extremely short time during which these critical regimes can be maintained. To this respect, we remind these results are obtained in the context of a simplified description of stellar rotation, and the reader should consider the remarks expressed in Sect.~\\ref{ssec_cautr}.   \\item Finally we recall that, according to a recent analysis by Baraffe et al. (2001), another possible mass-loss driving mechanism, namely the pulsation instability -- usually at work in VMO models with  ``normal'' chemical composition -- has been found of modest potential efficiency at zero metallicity, at least for masses $M \\la 500\\, M_{\\odot}$.  \\end{itemize}  By combining the supernova yields available in the literature with our predicted wind contributions, we also evaluate, in a simple way, the chemical enrichment produced by a primeval  population of (very) massive stars. It is found that a significant helium enrichment, $\\Delta Y \\sim 0.01$, may be reached by assuming that the primordial conditions in the Universe were such that only very massive stars could form, with typical masses of $\\approx 1000\\, M_{\\odot}$, or extending into the super-massive domain (up to $\\approx 10^{5} M_{\\odot}$).  In the former case the implied efficiency of gas-to-star conversion should be of the order of $10 \\%$, and the  first metal enrichment should be in the form of CNO elements (no iron-group elements). On the other side, under the assumption of the standard Salpeter IMF, the resulting $\\Delta Y$ is indeed negligible over the maximum allowed enrichment  in metals, $10^{-5} \\la \\Delta Z_{\\rm max} \\la 10^{-3}$. The corresponding star formation efficiency should be also be quite small, $\\approx 0.01-0.1 \\%$.  In any case, as already pointed out long ago by Bond et al. (1983) and Carr (1994), the interesting possibility arises  that the first generation of stars could  mask the true  (somewhat  lower) primordial abundance of helium as predicted by the standard Big Bang Nucleosynthesis (e.g. Olive et al. 1997). The question and its implications have been recently addressed by Salvaterra  \\& Ferrara (2002) on the base of the present results.   \\appendix"
    },
    "0212/astro-ph0212327_arXiv.txt": {
        "abstract": "{ We calculate the luminosity function of galaxies of the Early Data Release of the Sloan Digital Sky Survey (SDSS) and the Las Campanas Redshift Survey (LCRS). The luminosity function depends on redshift, density of the environment and is different for the Northern and Southern slice of SDSS. We use luminosity functions to derive the number and luminosity density fields of galaxies of the SDSS and LCRS surveys with a grid size of 1~$h^{-1}$ Mpc for flat cosmological models with $\\Omega _m=0.3$ and $\\Omega _\\Lambda =0.7.$ We investigate the properties of these density fields, their dependence on parameters of the luminosity function and selection effects.  We find that the luminosity function depends on the distance and the density of the environment. The last dependence is strong: in high-density regions brightest galaxies are more luminous than in low-density regions by a factor up to 5 (1.7 magnitudes). ",
        "introduction": "The study of the distribution of matter on large scales is usually based on the distribution of individual galaxies or clusters of galaxies. An alternative is to use the density field applying smoothing of galaxy or cluster distribution with a suitable kernel and smoothing length. This approach is customary in N-body simulations, where in each step a smoothed density field is evaluated. To the real cosmological data the smoothed density method has been applied in the study of the topology of the galaxy distribution by Gott et al. (\\cite{gmd86}). The IRAS redshift survey was used by Saunders et al. (\\cite{sfr91}) to calculate the density field up to a distance 140~$h^{-1}$\\thinspace Mpc. Recently Basilakos et al. (\\cite{bpr00}) applied the same method using the PSCz-IRAS redshift survey by Saunders et al. (\\cite{s00}). In all three studies a 3-dimensional spatial distribution was found. Due to the small volume density of galaxies with known redshifts a rather large smoothing length was used. This was sufficient to investigate topological properties of the galaxy distribution in the first case, and to detect superclusters of galaxies and voids in other cases.  In this paper, we shall calculate the number and luminosity density fields based upon the Early Data Release (EDR) of the Sloan Digital Sky Survey (SDSS) by Stoughton et al. (\\cite{s02}) and the Las Campanas Redshift Survey (LCRS) by Shectman et al. (\\cite{Shec96}). The SDSS Early Data Release consists of two slices of about 2.5 degrees thickness and $65-90$ degrees width, centred on celestial equator, LCRS consists of six slices of 1.5 degrees thickness and about 80 degrees width.  The number of galaxies observed per slice (over 10.000 in the SDSS slices and about 4,000 in LCRS slices) and their depth (almost 600~$h^{-1}$\\thinspace Mpc) are sufficient to calculate the 2-dimensional density fields with a high resolution (here $h$ is the Hubble constant in units of 100~km~s$^{-1}$~Mpc$^{-1}$). Using high-resolution number or luminosity density maps with smoothing scale of the order of 1~$ h^{-1}$\\thinspace Mpc\\ it is possible to find density enhancements in the field, which correspond to groups and clusters of galaxies. Using a larger smoothing length we can extract superclusters of galaxies as done by Basilakos et al. (\\cite{bpr00}). In calculating the density fields we can take into account most of the known selection effects, thus we hope that the density field approach gives additional information on the structure of the Universe on large scales. Clusters of galaxies from the SDSS were extracted previously by \\cite{2002PASJ...54..515G} using the cut and enhance method.  Loose groups of galaxies from LCRS were found by Tucker et al. (\\cite{Tucker00}). In accompanying papers by Einasto et al. (\\cite{e02a}, \\cite{e02b}, papers II and III, respectively) we use the density fields of SDSS and LCRS galaxies to derive catalogues of clusters and superclusters and to compare samples of density-field defined clusters and superclusters with clusters and superclusters found with conventional methods.   In the next section we describe the Early Data Release of the SDSS, and the LCRS samples of galaxies used. In section 3 we derive the luminosity function for the SDSS and LCRS galaxies using distances found for a cosmological model with dark matter and energy. In section 4 we calculate the density fields using SDSS and LCRS galaxy samples. We analyse our results in section 5, section 6 brings conclusions. ",
        "conclusions": "In this paper we have calculated the galaxy luminosity function for the SDSS EDR and LCRS samples and used it to construct the number and luminosity density fields (smoothed on $0.8\\ h^{-1}$ Mpc scale) assuming flat underlying cosmologies with $\\Omega _m=0.3$ and $\\Omega _\\Lambda =0.7.$ The analysis presented here is rather simple and serves as a first step in the study of the distribution of galaxies in SDSS and LCRS samples.  The principal conclusion from our study is: parameters of the galaxy luminosity function depend on the distance from the observer, density of the environment, they are different for the Northern and Southern slice.  The largest effect is the dependence on the density of the environment: in high-density regions brightest galaxies are more luminous than in low-density regions by a factor up to 5 (1.7 magnitudes).  Some of these effects suggest that it is not yet possible to find an universal set of parameters of the luminosity function valid for a fair sample of the Universe.  In other words, presently available samples are still too small to be considered as candidates for the fair sample."
    },
    "0212/astro-ph0212111_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:intro}  Most of the available information about matter in our Universe, and in our Galaxy in particular, comes indirectly from the collection of the electromagnetic radiation (from meter waves to $\\gamma$ rays) that was emitted or absorbed by this matter. A completely different information is provided by the cosmic ray nuclei, which constitute a genuine sample of galactic matter. Many different nuclei species are observed, in a wide range of energy, and with different origins. Some of them come unaltered from the sources (they are called {\\em primaries}\\/), others ({\\em secondaries}\\/) come from nuclear reactions between the primaries and the interstellar medium, or from the  disintegration of unstable species. Moreover, the trajectories of these nuclei from creation to detection are rather erratic, due to the influence of the galactic magnetic field on all the charged particles, and it is generally not possible to follow the direction of an incoming nucleus back to the source. If we were able to understand clearly the processes by which all these nuclei are produced, accelerated and propagated in the Galaxy, the wealth of data available now or in the near future would yield most valuable information about the matter content and magneto-hydrodynamical properties of our Galaxy. In principle, it would even be possible to discover some evidence for new physics (e.g. supersymmetry) or new objects (e.g. primordial black holes or stars made of antimatter) as they can give rise to the emission of charged  antinuclei and make an extra contribution to the observed cosmic ray fluxes. This review presents a summary of the work made in this direction by the authors from 1999 to 2002, in a {\\sc lapth-isn-iap} collaboration. As a first step, we tried to reach a {\\em quantitative}\\/ understanding of the propagation of cosmic ray nuclei in the energy range 100~MeV/nuc-100~GeV/nuc. More precisely, we described propagation with a diffusion model, in which the free parameters are adjusted to account for the available data on cosmic rays. This provides the regions of the parameter space allowed by the data. As a second step, we took advantage of this model to investigate several points concerning astrophysics and astroparticle physics.  During this study, various aspects  related to the ``standard\" or to more speculative processes were examined in detail. These may be summarized as follows: \\begin{itemize} \\item Standard cosmic rays \\begin{itemize} \\item Diffusion parameters from secondary-to-primary ratio \\item The flux of standard secondary antiprotons and antideuterons \\item Spatial origin \\item Radioactive species and the local bubble \\item Evolution of composition with energy \\end{itemize} \\item Exotic cosmic rays \\begin{itemize} \\item Baryonic Dark Matter \\item Antimatter \\item Supersymmetric particles \\item Primordial Black Holes \\end{itemize} \\end{itemize} The extraction of the diffusion parameters is the central goal since all conclusions follow from their values. ",
        "conclusions": "\\label{sec:concl}  A consistent framework to understand the propagation of CR nuclei in the energy range 100 MeV--100 GeV was presented in this paper. The observed fluxes of most species can be explained by assuming that once emitted from some sources located in the galactic disc, these nuclei undergo a diffusive propagation altered by escape through the boundaries, spallations, reacceleration, energy losses and galactic wind. The magnitude of these effects has been constrained using the B/C data, and the consistency of the model has been tested against the observed antiproton flux and by the study of radioactive species.  This well-tested model has then been used to study the propagation of cosmic rays of a more hypothetical origin, such as light antinuclei produced by SUSY galactic Dark Matter or Primordial Black Holes. In particular, limits on the abundance of primordial black holes in the galactic halo could be set.  There are several directions in which this work may be extended. First, the constraints on the propagation parameters could be refined by considering other species, stable or secondary. However, this approach is currently limited by the accuracy of the available data on cosmic ray fluxes and on the nuclear cross sections. Second, a specific study of the EC unstable species could provide valuable information about the  processes responsible for the acceleration of cosmic rays. Third, the existing limits on the SUSY induced antiproton signal can be bettered by using the constraints on the propagation parameters in a fully consistent way. Finally, the propagation code we use should ultimately be able to yield the flux of all cosmic ray species (including gamma rays, electrons and positrons) at every position in the Galaxy.   \\begin{figure}[hbt!] \\includegraphics[width=\\columnwidth]{organigramme_an.ps} \\caption{Schematic view of the subjects discussed in this paper. The stars indicate collaborations of {\\sc lapth} (Annecy-le-Vieux, France) members with {\\sc isn} ($\\star$, Grenoble, France), {\\sc iap} ($\\star \\star$, Paris, France) or {\\sc infn} ($\\star \\star \\star$, Turin, Italy). Dashed boxes represent future projects. The numbers in parenthesis represent the publications. The starting point (1) is the first use of the elaborate propagation model presented here \\cite{Maurin01}, (2) is \\cite{Donato02} (3) is \\cite{Donato01}, (4) is \\cite{barrau01}, (5) is \\cite{Maurin03}, (6) is \\cite{Maurin02}, (7) is \\cite{barrau02} and (8) is \\cite{Taillet01,Maurin04}. The pre-\\cite{Maurin01} works are labelled as (-1) for \\cite{Bottino_Salati} (-2) for \\cite{Orloff} and (-3) for \\cite{chardonnet96}.} \\label{fig:organigram} \\end{figure}"
    },
    "0212/astro-ph0212394_arXiv.txt": {
        "abstract": "We present a 40GHz (7.5 mm) raster scan image of a $3.6\\arcdeg\\times2\\arcdeg$ region centered on the low redshift ($z=0.055$) cluster of galaxies Abell 3667. The cluster was observed during the Antarctic winter of 1999 using the Corona instrument ($15.7\\arcmin$ FWHM beam) on the Viper Telescope at the South Pole. The Corona image of A3667 is one of the first direct ({\\it i.e.} rather than interferometer) thermal Sunyaev-Zel'dovich effect images of a low redshift cluster. The brightness temperature decrement at the X-ray centroid ($20^h 12^m 28.9^s, -56\\arcdeg 49\\arcmin 51\\arcsec$ J2000) was measured to be $\\Delta T_{\\rm CMB}=-154\\mu K$.  We have used the 40GHz map of A3667 in conjunction with a deep ROSAT PSPC (X-ray) image of the cluster, to make a measurement of the Hubble Constant. We find $H_0= 64^{+96}_{-30}$ km s$^{-1}$ Mpc$^{-1}$ (68\\% confidence interval).  Our $H_0$ calculation assumes that the cluster can be described using an isothermal, tri-axial ellipsoidal, $\\beta$-model and includes several new analysis techniques including an automated method to remove point sources from X-ray images with variable point spread functions, and an efficient method for determining the errors in multi-parameter maximum likelihood analyzes. The large errors on the $H_0$ measurement are primarily due to the statistical noise in the Corona image. We plan to increase the precision of our measurement by including additional clusters in our analysis and by increasing the sensitivity of the Viper SZE maps. ",
        "introduction": "\\label{sec:intro}  The Sunyaev-Zel'dovich effect (Sunyaev \\& Zel'dovich 1972, SZE hereafter) describes the inverse Compton scattering of cosmic microwave background (CMB) photons by energetic free electrons.  For instance, the random thermal motions of electrons trapped in the potential wells of clusters of galaxies result in a frequency dependent change in CMB intensity known as the thermal SZE. Similarly, the bulk peculiar motion of these electrons results in frequency independent change known as the kinetic SZE.  Observations of the thermal and kinetic SZE using either single dish telescopes ({\\it e.g.} Mason, Myers \\& Readhead 2001; Pointecouteau et al. 2001 \\& 2002; De Petris et al. 2002 or interferometric techniques ({\\it e.g.} Jones et al. 2001; Reese et al. 2001 \\& 2002; Udomprasert, Mason \\& Readhead 2001) have improved dramatically in recent years (see Birkinshaw 1999 for a recent review).  Measurements of the SZE have been used for a variety of scientific applications, including the estimation of the Hubble Constant (e.g. Mason, Myers, \\& Readhead 2001; Jones et al. 2001; Reese et al. 2002); determinations of the fraction, by mass, of baryons in the Universe (e.g. Grego et al. 2001); X-ray independent measurements of cluster temperatures (e.g. Pointecouteau et al. 2002); and constraints on cluster peculiar velocities (Holzapfel et al. 1997b, LaRoque et al. 2002).  The technique of interferometric SZE imaging is now well established. By contrast, only recently has it become possible to generate direct (rather than aperture synthesis) SZE images using single dish telescopes. To date, only two clusters, RX J1347-1145 \\& RX J2228+2037, both at $z\\simeq0.4$, have published direct SZE images (Pointecouteau et al. 1999, 2001 \\& 2002; Komatsu et al.  2001). We report here a third direct SZE image. This 40 GHz (7.5 mm) image of the ($z=0.055$, Sodre et al. 1992) Abell cluster (Abell, Corwin, \\& Olowin 1989) A3667 was made using the Viper telescope at the South Pole as part of the Viper Sunyaev-Zel'dovich Survey (VSZS). The VSZS aims to study a complete sample of southern clusters at radio/microwave, optical and X-ray wavelengths and it goals include the measurement of the Hubble Constant.  We chose to study A3667 as our initial VSZS target because, as one of the brightest X-ray ($L_{0.5-2.0 {\\rm keV}} = 4.1\\times10^{44}$ ergs $s^{-1}$, David et al. 1999) clusters in the REFLEX catalog (B\\\"ohringer et al. 2001), A3667 is expected to have a strong SZE signal. A3667 is also one of the best studied clusters in the sky. Supporting data at other wavelengths include: X-ray imaging and spectroscopy (ROSAT: Knopp, Henry \\& Briel 1996; ASCA: Markevitch, Sarazin, Vikhlinin 1999; White et al. 2000; BeppoSAX: Fusco-Femiano et al. 2001; Chandra: Vikhlinin, Markevitch \\& Murray 2001 a\\& b; XMM), optical data in the form of images, weak lensing mass maps and multi-object spectroscopy ({\\it e.g.}  Joffre et al. 2000; Katgert et al. 1998), and radio maps (R\\\"ottgering et al.  1997; Hunstead et al. 1999).  In addition, A3667 has been the focus of three-dimensional MDH/N-body numerical simulations (Roettiger, Burns \\& Stone 1999).   An outline of the paper is as follows. In section~\\ref{sec:viper-data} we review our observing strategy and data reduction methods. We also present the Corona map of the area around A3667. In section~\\ref{sec:Ho} we describe a joint fit, to an isothermal, tri-axial ellipsoidal $\\beta$-model, to the ROSAT Position Sensitive Proportional Counter (PSPC) and Corona images of A3667 and present an estimate of the Hubble Constant. In section~\\ref{sec:conclusions}, we present conclusions and discuss future plans for cluster observations with Viper. ",
        "conclusions": "\\label{sec:conclusions}  As the first stage of the Viper Sunyaev-Zel'dovich Survey (VSZS), we have presented a 40 GHz raster scan map of the region surrounding the $z=0.055$ cluster A3667. The cluster was observed during the Antarctic winter of 1999 using the Corona instrument ($15.7\\arcmin$ FWHM beam) on the Viper Telescope at the South Pole.  The Corona image of A3667 is one of the first Sunyaev-Zel'dovich effect (SZE) images of a low redshift cluster.  Currently, only the Viper telescope and the Cosmic Background Imager interferometer (Udomprasert, Mason \\& Readhead 2001) are able to make SZE images of low redshift clusters. One of the advantages of low redshift observations of the SZE is that complimentary data at other wavelengths, for example those necessary to make Hubble Constant estimates, are comparatively easy to obtain. The Corona image of A3367 is also one of the first direct ({\\it i.e.}  rather than interferometer) SZE images; only two other clusters have published direct SZE images.  We have used the 40GHz map of A3667 in conjunction with a deep ROSAT PSPC (X-ray) image to make a maximum likelihood fit to an isothermal tri-axial ellipsoidal $\\beta$ model. We have determined 68\\% confidence regions for nine of the eleven free parameters in the model using a modified Likelihood Ratio Test. Our analysis method includes new approaches to point source detection in X-ray images with varying point spread functions and also to error estimation in multi-dimensional parameter space. These innovations will be applied during the analysis of other VSZS clusters, but are also relevant to other SZE experiments. Our measurement of the Hubble Constant, $H_0= 64^{+96}_{-30}$ km s$^{-1}$ Mpc$^{-1}$, is in good agreement with previous, SZE determined, measurements. The error on this measurement is large, but not atypical for single dish experiments, and is being driven by instrument noise in the Corona map. Other parameters in the fit have been determined to much higher precision (1-3\\%).  The Corona instrument was retired at the end of 2000 and replaced with the ACBAR instrument (Runyan et al. 2002). ACBAR offers both improved spatial resolution ($5\\arcmin$ FWHM beam) compared to Corona and simultaneous multi-wavelength (150, 220, 280 GHz) capabilities.  Using ACBAR, in combination with data from CTIO, XMM-Newton and Chandra, we have continued the VSZS and are working toward completion of the SZE, X-ray and weak lensing observations of a complete, luminosity limited, sample of $\\simeq10$ clusters.  These observations will improve our $H_0$ measurement in several ways. First, long duration ACBAR observations produce more sensitive maps than was possible with Corona (see Kuo et al. 2002). Second, we can minimize the contribution of primary CMB anisotropies in our SZE images by combining data at the three ACBAR observing frequencies (Gomez et al. 2002). This feature becomes important when the instrument noise gets down to the level of the CMB noise and it is noteworthy that most other SZE experiments only operate at a single frequency; Diabolo (Pointecouteau 1999, 2001 \\& 2002), BOLOCAM (Mauskopf et al. 2000b) and SuZIE (Holzapfel 1997 a\\&b) are exceptions, but these are not optimized to image low redshift clusters. Third, combining observations of several clusters drawn from a statistically complete sample reduces both random errors and systematic errors associated with orientation bias ({\\it e.g.} Grainger 2001). Finally, the improved spatial resolution of ACBAR over Corona, means that we will be able to use independent beams to probe different lines of the sight through each cluster. We only used data from the central beam of the A3667 Corona image for the analysis above, but by using multiple lines of sight we can reduce the error on the overall measurement. Together, these improvements will allow us to measure the Hubble constant with significantly higher precision than was possible with the Corona observation of A3667 presented herein.  \\paragraph"
    },
    "0212/astro-ph0212441_arXiv.txt": {
        "abstract": "We are modeling the spectra of dwarf galaxies from infrared to submillimeter wavelengths to understand the nature of the various dust components in low-metallicity environments, which may be comparable to the ISM of galaxies in their early evolutionary state. The overall nature of the dust in these environments appears to differ from those of higher metallicity starbursting systems. Here, we present a study of one of our sample of dwarf galaxies, NGC~1569, which is a nearby, well-studied starbursting dwarf. Using ISOCAM, IRAS, ISOPHOT and SCUBA data with the D\\'esert et al (1990) model, we find consistency with little contribution from PAHs and Very Small Grains and a relative abundance of bigger colder grains, which dominate the FIR and submillimeter wavelengths. We are compelled to use 4 dust components, adding a very cold dust component, to reproduce the submillimetre excess of our observations. ",
        "introduction": "Dwarf galaxies in our local universe are ideal laboratories for studying the interplay between the ISM and star formation in low-metallicity environments (Hunter \\& Gallagher 1989). They are at relatively early epochs of their chemical evolution, possibly resembling distant protogalaxies in their early stages of star formation. Although only a few metal-poor galaxies have been observed in the MIR using ISO, the characteristics of the MIR dust components appear to differ remarkably from those of normal-metallicity starbursts (Madden 2000). Whether this is an abundance or composition effect is not yet clear.  For this reason, we are studying their detailed luminosity budget by modeling their spectral energy distributions (SEDs) from optical to millimeter wavelengths, thereby constructing templates to study conditions in primordial galaxies to help to constrain galaxy evolution models.  The galaxy we present here is NGC~1569. It is a HI-rich, metal-poor ($Z=0.25\\;\\rm Z_\\odot$) dwarf irregular galaxy which lies near the galactic plane at a distance of $(2.2 \\pm 0.6)\\;\\rm Mpc$ (Israel 1988). NGC~1569 is presently in the aftermath of a massive burst of star formation (Israel 1988, Israel \\& De Bruyn 1988, Waller 1991) and exhibits very compact HII regions and 2 super-star-clusters (Hunter et al 2000). ",
        "conclusions": "The dust masses deduced from our modeling are $m_{\\rm PAH}\\sim 3 \\msol$, $m_{\\rm VSG}\\sim 600 \\msol$, $m_{\\rm BG}\\sim 2.5\\, 10^5 \\msol$, $m_{\\rm VCG}\\sim 2.7\\, 10^5 \\msol$ and the gas-to-dust mass ratio is $\\sim 400$. % The dust mass deduce from the extinction is $\\sim 2.2\\, 10^5 \\msol$ which is the same order of magnitude as the mass deduced from the model. Our results seem to show that the dust mass in NGC~1569 is mainly concentrated in cold dust: big grains and very cold grains which both have roughly the same mass abundance. The PAHs as well as the VSGs are very sparse, while the SED is dominated by bigger grains."
    },
    "0212/astro-ph0212507_arXiv.txt": {
        "abstract": "We present interplanetary network localization, spectral, and time history information for 7 episodes of exceptionally intense gamma-ray emission from Cygnus X-1.  The outbursts occurred between 1995 and 2003, with durations up to $\\thicksim$28000 seconds.  The observed 15 - 300 keV peak fluxes and fluences reached $\\rm3 \\times 10^{-7}\\, erg \\, cm^{-2} s^{-1} \\, and \\, 8 \\times 10^{-4}\\, erg \\,cm^{-2}$ respectively.  By combining the triangulations of these outbursts we derive an $\\thicksim$ 1700 square arcminute (3 $\\sigma$) error ellipse which contains Cygnus X-1 and no other known high energy sources. The outbursts reported here occurred both when Cyg X-1 was in the hard state as well as in the soft one, and at various orbital phases.  The spectral data indicate that these outbursts display the same parameters as those of the underlying hard and soft states, suggesting that they represent another manifestation of these states. ",
        "introduction": "The X-ray source Cygnus X-1 was discovered by Bowyer et al (1965), and its optical counterpart, a spectroscopic binary at a distance of $\\sim$ 2 kpc with a 5.6 day period (HDE226868), was identified by Webster and Murdin (1972) and Bolton (1972).  The primary is a supergiant, and the mass of the secondary is at least $\\rm 7 M_{\\sun}$ (Gies \\& Bolton 1986), making it a black hole candidate and a high mass X-ray binary.  The X-ray source, which has been observed from keV to MeV energies, is thought to be powered by accretion (e.g. Petterson 1978).  Soft X-radiation may be produced in an accretion disk close to the black hole, and hard X-rays by inverse Compton scattering in a separate hot plasma; see Liang and Nolan (1984) for a review.  Recently, a relativistic jet has been detected in the radio (Stirling et al. 2001), and Romero et al. (2002) have suggested that Cygnus X-1 is a microblazar, that is, that we are observing it close to the axis of the jet.  Cygnus X-1 is known to exhibit two states of X-ray emission, which are thought to depend on the accretion rate (Esin et al. 1998).  The first, more common one, is the hard state, in which the 20 - 200 keV flux is roughly one Crab, and the $<$10 keV emission is about 0.5 Crab.  The second is the soft state, in which the low energy X-ray flux increases by a factor of 2 - 4, while the high energy flux decreases by a factor of about 2.  Transitions between the two states and flaring have been documented extensively (Ling et al. 1997; Zhang et al. 1997; Cui et al. 2002; McConnell et al. 2002), and both states exhibit variability of a factor of two or so on all timescales.  In this paper we report on seven episodes of long, intense gamma-ray emission from this source; we refer to them as \\it outbursts \\rm (and identify them by their dates), to distinguish them from the shorter, less intense flaring which has been documented previously.  One of these outbursts, 950325, was found by Mazets et al. (1995) and initially thought to be a gamma-ray burst; subsequently, however, its source was found to be Cygnus X-1 by M. Briggs (private communication, 1995).  Later, Stern et al. (2001) called attention to five outbursts on 1999 April 19-21.  During the two brightest events, the hard X-ray flux increased by over an order of magnitude for $\\sim$ 1000 s. Observations by \\it Ulysses \\rm and Konus - \\it Wind \\rm of still three more outbursts in 2002 - 2003 have prompted us to reanalyze all the available data in order to obtain better energy spectra and a more precise source location as well. ",
        "conclusions": "The outbursts reported here have durations comparable to the period of a low Earth-orbiting spacecraft, making them difficult to detect and follow from experiments aboard such spacecraft, but relatively easy for experiments which are far from Earth and do not undergo occultation and orbital background variations.  Indeed, our observations of 990421A show that it continues well beyond the point at which it was Earth-occulted to BATSE, and our observations of 990421B indicate that it commenced several hundred seconds before it rose on BATSE (Stern et al. 2001).  These observations point to a new use for the IPN, namely tracking long, intense flares from Galactic transients.  The data from a single experiment which has little or no directional and/or spectral information are easily confused with solar X-ray and particle events, which explains why it has taken so long to determine the origin of some of the outbursts presented here.  However, confirmation by a second spacecraft solves this problem and gives an annulus of position which in many cases may be sufficient to determine the origin of the emission.   For several outbursts, the BATSE observations alone yield source directions which are only accurate to 17$\\degr$ ,  leaving the possibility that the source could have been Cygnus X-3.  However, relatively small error ellipses may be obtained from multiple observations, and our localizations rule this possibility out conclusively.  The histories of Cygnus X-1 and GRBs have been curiously intertwined over the decades.  In the early days of X-ray astronomy, the use of interplanetary spacecraft was suggested to localize sources with rapid time variations such as Cyg X-1 (Giacconi 1972).  Although, to our knowledge, such measurements were never carried out on persistent X-ray sources, today the technique is the basis of the IPN, and the present observations demonstrate the feasibility of this suggestion.  Later, Mason et al. (1997) pointed out that bursts from Cyg X-1 could appear in the BATSE database; BATSE has indeed triggered many times on this source.  Some of the long outbursts presented here have time histories which, if compressed in time, would be virtually impossible to distinguish from those of GRBs, and their spectra are hard.  These similarities suggest that accretion onto a black hole, which is believed to power both Cyg X-1 and GRBs albeit under very different circumstances, may manifest itself in similar ways in very different settings."
    },
    "0212/astro-ph0212076_arXiv.txt": {
        "abstract": "The evolution of the magnetorotational instability (MRI) during the transition from outburst to quiescence in a dwarf nova disk is investigated using three-dimensional MHD simulations. The shearing box approximation is adopted for the analysis, so that the efficiency of angular momentum transport is studied in a small local patch of the disk: this is usually referred as to a one-zone model. To take account of the low ionization fraction of the disk, the induction equation includes both ohmic dissipation and the Hall effect. We induce a transition from outburst to quiescence by an instantaneous decrease of the temperature. The evolution of the MRI during the transition is found to be very sensitive to the temperature of the quiescent disk. As long as the temperature is higher than a critical value of about 2000 K, MHD turbulence and angular momentum transport is sustained by the MRI. However, MHD turbulence dies away within an orbital time if the temperature falls below this critical value. In this case, the stress drops off by more than 2 orders of magnitude, and is dominated by the Reynolds stress associated with the remnant motions from the outburst. The critical temperature depends slightly on the distance from the central star and the local density of the disk. ",
        "introduction": "Dwarf novae are close binary systems that consist of a white dwarf and a Roche-lobe-filling secondary star. The matter overflowing from the secondary star forms an accretion disk around the white dwarf. Dwarf nova systems show repetitive outbursts in which the luminosity of the disks increase rapidly. The duration of the outbursts ranges from a few to 20 days, and the recurrence time is about 20 -- 300 days. The disk instability model (Osaki 1974), in which the evolution of the disk is regulated by changes in the rate of angular momentum transport, is generally accepted as the explanation of the outbursts (Cannizzo 1993; and references therein). Since ordinary molecular viscosity is too low to explain the evolutionary timescale of the disks, an anomalous stress associated with turbulence and characterized by the $\\alpha$ parameter (Shakura \\& Sunyaev 1973), is thought to be the source of angular momentum transport. Detailed comparisons between observed light curves of dwarf novae and the theoretical disk models suggest that the viscous parameter $\\alpha$ must vary between the hot (outburst) and cold (quiescent) state, and the amplitude of $\\alpha_{\\rm hot}$ and $\\alpha_{\\rm cold}$ is typically of the order of 0.1 and 0.01, respectively. However, while the origin of the anomalous stress has become better understood in recent years (see below), the reasons for the difference between $\\alpha_{\\rm hot}$ and $\\alpha_{\\rm cold}$ are still unclear.  The magnetorotational instability (MRI; Balbus \\& Hawley 1991) is the most promising source of the anomalous stress. MHD turbulence driven by the MRI can transport angular momentum by the Maxwell (magnetic) stress. During the hot state, the disk gas is fully ionized so that the ideal MHD approximation is appropriate. Numerical simulations that adopt ideal MHD have shown that the Maxwell stress caused by the MRI can give $\\alpha \\sim 0.1$ to 0.01, where the saturation level is determined by the geometry and strength of the magnetic field (e.g., Hawley, Gammie, \\& Balbus 1995; 1996). Therefore, MHD turbulence can be the primary mechanism of angular momentum transport at least during outbursts, that is it can account for $\\alpha_{\\rm hot}$.  During the cold state, on the other hand, the gas is only weakly ionized. Gammie \\& Menou (1998) have pointed out that ohmic dissipation could modify the nature of the MHD turbulence at quiescence, and this could make the difference between $\\alpha_{\\rm hot}$ and $\\alpha_{\\rm cold}$. Calculations of the ionization fraction in dwarf nova disks reveal that the Hall effect as well as ohmic dissipation must be considered if the temperature is $T \\lesssim 2000$ K (Sano \\& Stone 2002a). At low temperatures, ohmic dissipation can suppress the growth of the MRI (Jin 1996; Sano \\& Miyama 1999). If turbulence dies away completely as a result of ohmic dissipation, the Maxwell stress cannot provide the required anomalous stress during quiescence (Menou 2000). Thus, it is important to calculate numerically the amplitude of the Maxwell stress at the cold state.  In this paper, we examine the behavior of MHD turbulence during the transition from an outburst to quiescence using local three-dimensional MHD simulations.  These simulations include the dominant non-ideal MHD effects as determined from a self-consistent calculation of the ionization state in the disk.  Since we adopt a local approximation to study the turbulent stresses (usually referred to as a one-zone model, e.g., Mineshige \\& Osaki 1983; Cannizzo \\& Wheeler 1984), our calculations do not include global effects (such as spiral shocks in the disk, Sawada, Matsuda, \\& Hachisu 1986) which may contribute to angular momentum transport.  The plan of this paper is as follows.  In \\S~2, we calculate the ionization fraction in dwarf nova disks by solving the Saha equation to estimate the magnitude of nonideal MHD effects at quiescence.  Our numerical method is described in \\S~3.  The results of numerical simulations which examine the dependence of the stress on the geometry of magnetic field and the decay timescale of turbulence are presented in \\S~4.  In \\S~5, we discuss the activity of MHD turbulence and the source of angular momentum transport during quiescence. ",
        "conclusions": "\\subsection{The Critical Temperature and Transition Timescale}  When the magnetic Reynolds number falls to unity, MHD turbulence is suppressed by ohmic dissipation independent of the amplitude of the Hall term (Sano \\& Stone 2002a,b). The critical temperature corresponding to $Re_{M} = 1$ is about 2000 K for our model. As seen from equation~(\\ref{eqn:rem}), however, the critical temperature depends on the neutral density $n_n$ and the distance from the central star $r$. Figure~\\ref{fig:tcrit} shows the critical temperature $T_{\\rm crit}$ as a function of $r$ for the cases of the neutral density $n_n = 10^{17}$, $10^{18}$, and $10^{19}$ cm$^{-3}$. We assume the Alfv{\\'e}n speed $v_{A} = 10^5$ cm s$^{-1}$ and the mass of the central star $M = M_{\\odot}$. At a given temperature, the electron abundance decreases as the neutral density increases, so that the critical temperature is slightly higher when the density is higher. The MRI wavelength ($\\sim v_{\\rm A} / \\Omega$) is shorter at a smaller radius, because the angular velocity is larger. Therefore, the MRI can be suppressed with a smaller amount of diffusivity $\\eta$ in the inner regions of the disk. The critical temperature is about 2500 and 1500 K at $r = 10^9$ and $10^{11}$ cm, respectively. Since the magnetic Reynolds number, or the electron abundance, has a steep dependence on the temperature, the range of the critical temperature is quite narrow and the difference is at most a factor of 2 over the entire disk.  In our numerical analysis, we assume a fixed temperature during quiescence. Thus the electron abundance at quiescence is assumed to be constant. In a real disk, the critical temperature could be a little higher due to a runaway decay of MHD turbulence (Menou 2000). The decrease of the temperature makes both the electron abundance lower and nonideal MHD effects more important. The turbulent stress is therefore suppressed further, so that the temperature falls even more. However, a very weak dependence of the stress on the temperature can be seen in our models when $T \\gtrsim 3000$ K. This indicates that the runaway temperature should be less then 3000 K. Although a self-consistent study of the evolution of the temperature and electron abundance is needed to obtain a real $T_{\\rm crit}$, Figure~\\ref{fig:tcrit} may give a fairly good estimate of the critical temperature in dwarf nova disks.  We also assume the transition to quiescence occurs instantaneously. In actual systems, however, this may proceed at the thermal timescale $t_{\\rm th}$, where $t_{\\rm th}/t_{\\rm rot} \\sim 1 / \\alpha \\gtrsim 100$ (Frank, King, \\& Raine 1992). Suppose the temperature decreases gradually. As long as the temperature is higher than the critical value, we have shown that the accretion stress is nearly constant. But, once the temperature drops below this critical value, the MRI is suppressed,  the turbulence decays within an orbit, and the $\\alpha$ parameter drops by more than 2 orders of magnitude. Therefore, even though the decrease of the temperature is gradual, the transition timescale of the stress may be very much shorter.  \\subsection{Onset of the Next Outburst}  Dwarf nova systems undergo recurrent outbursts. The matter from the secondary star accumulates in the disk during quiescence. When the surface density exceeds a critical value at some region in the disk, the thermal instability sets in and the temperature goes up. Here we examine the behavior of the MRI at the transition to the hot state.  Figure~\\ref{fig:w-t100-z} shows the time evolution of the Maxwell and Reynolds stress for model ZC until 100 orbits. After 100 orbits, the induction equation for the ideal MHD is used in order to imitate the onset of the next outburst. As seen from the figure, the growth of the MRI starts immediately at the transition to the outburst and the magnetic field is amplified exponentially. In a few orbits, the Maxwell stress reaches the same saturation level as that before 50 orbits. The key to this evolution is the existence of a net vertical flux. During quiescence (from 50 to 100 orbits), MHD turbulence is suppressed, but the magnetic flux within the shearing box is conserved throughout the evolution. Thus the field geometry at 100 orbits is a nearly uniform vertical field. When the nonideal MHD effects are turned off, the linear growth of the MRI starts and MHD turbulence is initiated.  This evolution may be different if the system has no net magnetic flux. For example, if the magnetic field is {\\em completely} dissipated during quiescence (e.g., model SD), MHD-driven activity cannot occur during the next outburst. However, if enough magnetic field remains for there to be unstable modes even in a small region of the disk, MHD turbulence develops in that region and finally may spread over a large part of the disk (Hawley \\& Balbus 1991). Therefore, it is quite important to understand the global structure of the magnetic field during quiescence, including the magnetosphere of the white dwarf and the magnetic field supplied by accreting material from the secondary star (Meyer \\& Meyer-Hofmeister 1999).  \\subsection{Beyond the Local Model}  Since we focus on the local behavior of MHD turbulence in this paper, neither global structural effects nor global instabilities have been considered. In real disks, global effects such as spiral shocks (Sawada et al. 1986) may also transport angular momentum.  From our numerical results, the dependence of the $\\alpha$ parameter associated with local turbulence on temperature in a dwarf nova disk can be summarized as follows.  The contribution of the MRI to the $\\alpha$ parameter is significant during outburst.  When the temperature is higher than the critical value $T_{\\rm crit} \\sim 2000$~ K, the accretion stress is dominated by the Maxwell stress, with $\\alpha \\sim 0.1$ and 0.01 for the cases with and without net vertical flux, respectively.  At quiescence, on the other hand, $\\alpha$ is very sensitive to the temperature of the disk.  If $T \\lesssim T_{\\rm crit}$, the stress is more than 2 orders of magnitude smaller than that during outburst.  There are several reasons why global effects may influence these results.  Firstly, the presence of a net vertical flux can affect the amplitude of the turbulence stress during outburst, that is $\\alpha_{\\rm hot}$.  In agree with the previous works using the local shearing box (e.g., Hawley et al. 1995; 1996), our simulations have shown that an sufficient amount of angular momentum transport in the hot state requires the existence of a net vertical field.  On the other hand, recent global simulations obtained a higher stress ($\\alpha \\sim 0.1$) even without a net magnetic flux (Stone \\& Pringle 2001; Hawley, Balbus, \\& Stone 2001).  There is no observational constraint yet in the structure and strength of the magnetic field in dwarf nova disks. Thus, theoretical modeling of the global structure of the magnetic field is quite important for a more detailed analysis of $\\alpha$ at the hot state. Secondly, the ionization fraction of the disk gas, and therefore the importance of non-ideal MHD effects, is extremely sensitive to temperature.  Thus global models in which the radial and vertical temperature profiles in the disk are calculated self-consistently with the local heating rate due to turbulent stresses and cooling rate due to radiation are required. Furthermore, at the surface of the disk, the radiation from a hot white dwarf could be important source of heating and ionization (Hameury, Lasota, \\& Dubus 1999; Menou 2002). Finally, even when Ohmic dissipation completely suppresses the MRI, small amplitude motions remain in the disk and can give a non-zero Reynolds stress.  In the local model, these motions are damped on a viscous timescale, which may be very long.  However, they may be affected by global density waves in the disk or tidally induced spiral shocks. Thus, global MHD simulations are an important next step for understanding the dynamics of dwarf nova disks."
    },
    "0212/astro-ph0212240_arXiv.txt": {
        "abstract": "We present optical and near--infrared Keck spectroscopy of CXOHDFN J123635.6$+$621424 (hereafter HDFX28), a hard X--ray source at a redshift of $z = 2.011$ in the flanking fields of the Hubble Deep Field--North (HDF--N). HDFX28 is a red source (${\\cal R} - K_s = 4.74$) with extended steep--spectrum ($\\alpha^{\\mbox{\\tiny 8.4 GHz}}_{\\mbox{\\tiny 1.4 GHz}} > 0.87$) microjansky radio emission and significant emission (441 $\\mu$Jy) at 15 $\\mu$m. Accordingly, initial investigations prompted the interpretation that HDFX28 is powered by star formation.  Subsequent {\\it Chandra} imaging, however, revealed hard ($\\Gamma = 0.30$) X--ray emission indicative of absorbed AGN activity, implying that HDFX28 is an obscured, Type II AGN.  The optical and near--infrared spectra presented herein corroborate this result;  the near--infrared emission lines cannot be powered by star formation alone, and the optical emission lines indicate a buried AGN. HDFX28 is identified with a face--on, moderately late--type spiral galaxy. Multi--wavelength morphological studies of the HDF--N have heretofore revealed no galaxies with any kind of recognizable spiral structure beyond $z > 2$. We present a quantitative analysis of the morphology of HDFX28, and we find the measures of central concentration and asymmetry to be indeed consistent with those expected for a rare high--redshift spiral galaxy. ",
        "introduction": "\\label{introduction}  The origin of the X--ray background (XRB) remains an enduring puzzle for X--ray astronomy.  Great progress has been made during the last four years: {\\it ROSAT} surveys successfully resolved $\\sim$80\\% of the soft XRB (0.5--2 keV) into discrete sources \\citep[e.g.][]{hasinger98} and current work with the {\\it Chandra X--ray Observatory} (hereafter, {\\it Chandra})  has successfully resolved a similar fraction of the hard XRB \\citep[2--8 keV;][]{brandt01, giacconi01, giacconi02, hornschemeier01, rosati02}.  Nonetheless, a coherent understanding of the physical and evolutionary properties of the sources which comprise the XRB is only just now emerging. Although much of this population appears to be the nuclei of otherwise normal bright galaxies ($I < 23.5$) or typical active galactic nuclei (AGNs), a significant fraction of the discrete sources is optically faint ($I > 23.5$), and therefore not easily identified \\citep[e.g.][]{alexander01, barger02}.  {\\it Type II quasars}, for instance, are thought to be AGNs viewed edge--on through an obscuring torus \\citep{antonucci93} and are deemed an essential component of the XRB--producing population \\citep{moran01}. However, few well--studied examples of such systems are known at high redshift, and owing in part to their lack of relativistic brightening, they are not easily identified in shallow, large--area surveys \\citep{norman02, stern02q}.  Type II quasars represent just one of several diverse classes of objects emerging in follow--up work to deep {\\it Chandra} fields \\citep[e.g.][]{hornschemeier01, schreier01, stern02}. On the extra--galactic side, we find X-ray--loud composite galaxies typified by starburst or early--type optical spectra which bear no signature of their buried AGN \\citep[e.g.][]{moran96, levenson01, stern02}. Additionally, we find X--ray sources whose optical counterparts belong to the class of faint, extremely red objects (EROs), the nature of which has remained uncertain owing to the difficulty in spectroscopic follow--up \\citep[e.g.][]{alexander02, elston88, elston89, hu94, graham96, liu00, hornschemeier01, stern02}. On the Galactic side, we find late--type dwarfs emitting soft X--rays originating in chromospheric activity \\citep[e.g.][]{hornschemeier01}, and very low mass binary systems emitting hard X--rays driven by accretion \\citep[e.g.][]{stern02}.  Amidst the emergence of this menagerie of objects, optical and near--infrared spectroscopic follow--up has become increasingly vital not only to identifying the source population of the XRB, but also to elucidating the physics of X--ray sources in general, and to delineating their evolution with redshift.  One critical facet of this endeavor is simply to distinguish between objects powered by mass accretion onto supermassive black holes (quasars and other AGNs) and those powered by nuclear fusion in stars (normal and starburst galaxies).  To this end, we present optical and near--infrared spectra of CXOHDFN J123635.6$+$621424 (hereafter HDFX28), a hard X--ray source identified with a face--on spiral galaxy at redshift $z = 2.011$ (Figure~\\ref{flank}). HDFX28 is fortuitously located in the Hubble Deep Field--North inner west (HDF--N IW) flanking field, and was therefore subject to a vast array of follow--up imaging. As such, HDFX28 was initially identified as an extended microjansky radio source with a comparatively steep spectral index ($S_{\\nu} \\propto \\nu^{-\\alpha}$; $\\alpha^{\\mbox{\\tiny 8.4 GHz}}_{\\mbox{\\tiny 1.4 GHz}} > 0.87$).  Together with its detection by the {\\em Infrared Space Observatory} Camera (ISOCAM) and its pronounced optical spatial extent ($\\sim 1\\farcs6$), the radio data for HDFX28 prompted an initial interpretation as a galaxy powered by star formation \\citep[e.g.][]{richards00}. However, as we discuss below, the detection of HDFX28 as a hard X--ray source in the deep {\\it Chandra} survey of the HDF--N \\citep{hornschemeier01, brandt01}, corroborated by the spectroscopy presented herein, demonstrates that this galaxy in fact harbors an obscured, Type II AGN.  In addition to confirming its AGN status, the spectroscopy of HDFX28 indicates a surprisingly high redshift for an object with identifiable spiral structure.  \\citet{dickinson00} summarizes the results of morphological studies of the HDF--N by reporting a total lack of even plausible candidates for spiral galaxies at $z>2$.  Prompted by this lack of precedent for high--redshift spirals, we present a quantitative study of central concentration and asymmetry in HDFX28 based on the scheme devised by \\citet{abraham96} for the analysis of large CCD imaging surveys. With the application of a modest morphological $k$--correction, we find HDFX28 to have morphological parameters consistent with those derived from catalogs of both artificially redshifted nearby spirals, as well as catalogs of {\\it HST} imaging of spirals out to $z \\sim 1$ \\citep{abraham96}.  In short, HDFX28 is intriguing both for its membership in the emerging class of X--ray--selected Type II AGN, and for possessing a morphology which is unprecedented at its redshift. We describe the optical and near--infrared spectroscopy of HDFX28 in section \\S \\ref{observations}, and we present the results of the spectroscopy and the classification of the source as an obscured, Type II AGN in \\S \\ref{as_an_agn}.  We report on our quantitative analysis of its morphology in \\S \\ref{as_a_spiral}, and we summarize our results in \\S \\ref{conclusion}. Throughout this paper we adopt the currently favored $\\Lambda$--cosmology of $\\Omega_{\\mbox{\\tiny M}} = 0.35$ and $\\Omega_\\Lambda = 0.65$, with $H_0 = 65$ km s$^{-1}$ Mpc$^{-1}$ \\citep[e.g.][]{riess01}.  At $z=2.011$, such a universe is 3.22 Gyr old, the lookback time is 76.9\\% of the total age of the Universe, and an angular size of 1\\farcs0 corresponds to 8.66 kpc.  \\bigskip \\bigskip \\bigskip ",
        "conclusions": "\\label{conclusion}  We have reported on two aspects of the high--redshift, hard X--ray emitting spiral galaxy HDFX28: (1) its classification as a Type II AGN, a population recently attracting renewed interest due to deep X--ray surveys, and for which few {\\it HST} images are available, and (2) its unprecedented redshift for a galaxy with spiral morphology. As for HDFX28 as a Type II AGN, the canonical wisdom regarding weak, extended radio sources with spectral indices steeper than $\\alpha^{\\mbox{\\tiny 8.4 GHz}}_{\\mbox{\\tiny 1.4 GHz}} > 0.5$ dictates that such sources are driven by star formation. Nonetheless, the combined weight of evidence from X--ray, optical, and near--infrared observations of HDFX28 indicates the presence of obscured AGN activity. It is instructive to note that when re--interpreted in light of the spectroscopic redshift, even the mid--infrared data for HDFX28 corroborates this result. At $z = 2.011$, the ISOCAM LW3 filter samples rest wavelengths spanning only 4 $\\mu$m to 5 $\\mu$m. Here, the contribution to the mid--IR spectral energy distribution made by UIB emission and by dust at 200 K is severely attenuated \\citep[see][Figure 1]{aussel99}.  Hence, the ISOCAM detection of this source is far more plausibly explained by the hot, $\\sim 10^3$ K dust found in the central region of an AGN \\citep[e.g.\\ see][]{aussel98} than it is by star formation alone.  As to the precise nature of the central engine in HDFX28, we conclude from the comparatively narrow emission lines in the spectroscopy and from the heavy obscuration evident in the X--ray data that HDFX28 is far more like an obscured Type II system than an unobscured Type I system. Though this conclusion is slightly at odds with the presence of weak, broad \\hal\\ emission, all remaining aspects of the source are entirely consonant with observations of other Type II AGN at moderate--to--high redshifts \\citep[e.g.][]{kleinmann88, norman02, stern02q} and with HzRGs \\citep[e.g.][]{larkin00, mccarthy93, stern99hzrg, vernet01}. \\citet{norman02} describe a very similar situation in which their source CDF--S 202 shows both the narrow ($\\sim 1000$ km s$^{-1}$) emission lines in its optical spectrum and the heavy obscuration in its X--ray emission typical of a Type II system, but also shows emission line flux ratios intermediate between Type I and Type II systems.  As noted by \\citet{stern02q}, it is conceivable that longer--wavelength spectra of CDF--S 202 and other sources like it would also reveal broad \\hal, though they in every other way give evidence of the heavy obscuration considered to be emblematic of Type II AGN.  Separately, the spectroscopy presented herein shows HDFX28 to be at an unprecedented redshift for a galaxy with identifiably spiral structure. Nevertheless, with the application of a modest morphological $k$--correction, our quantitative analysis of its central concentration and asymmetry is consistent with the interpretation that HDFX28 is a rare example of a high--redshift spiral galaxy.  Owing to its proximity to the HDF--N, HDFX28 will be subject to deep, space--based $B$, $V$, $i$, and $z$ imaging with the Advanced Camera for Surveys as part of the upcoming GOODS {\\it HST} Treasury Program (M.\\ Giavalisco, PI), as well as to infrared imaging at $\\lambda > 3$ $\\mu$m with the Infrared Array Camera as part of the GOODS {\\it SIRTF} Legacy project (M.\\ Dickinson, PI).  At a minimum, the availability of multi--wavelength imaging will provide a powerful additional lever arm on the issue of the morphology of HDFX28 \\citep[e.g.][]{conselice97,conselice00}. As such, we eagerly look forward to these expansive datasets."
    },
    "0212/astro-ph0212289_arXiv.txt": {
        "abstract": "We report the first measurements of anisotropy in the cosmic microwave background (CMB) radiation with the Arcminute Cosmology Bolometer Array Receiver ({\\sc Acbar}). The instrument was installed on the $2.1\\,$m Viper telescope at the South Pole in January 2001; the data presented here are the product of observations up to and including July 2002. The two deep fields presented here, have had offsets removed by subtracting lead and trail observations and cover approximately $24\\,{\\rm deg}^2$ of sky selected for low dust contrast. These results represent the highest signal to noise observations of CMB anisotropy to date; in the deepest $150\\,$GHz band map, we reached an RMS of $\\sim8.0\\,\\mu$K per $5^{\\prime}$ beam. The 3 degree extent of the maps, and small beamsize of the experiment allow the measurement of the CMB anisotropy power spectrum over the range $\\ell = 150-3000$ with resolution of $\\Delta \\ell=150$. The contributions of galactic dust and radio sources to the observed anisotropy are negligible and are removed in the analysis. The resulting power spectrum is found to be consistent with the primary anisotropy expected in a concordance $\\Lambda$CDM Universe. ",
        "introduction": "\\ln{\\cal L}({\\bf q}) = \\ln{\\cal L}(\\overline{{\\bf q}})-\\frac{1}{2} \\sum\\limits_{B}\\frac{(Z_B - \\overline{Z}_B)^2}{\\sigma_B^2}e^{2{\\overline Z}_B}, \\end{equation} where the offset log-normal parameters ${\\bf Z}$ are defined as \\begin{equation}\\label{lognorm} Z_B = \\ln{(q_B + x_B)}. \\end{equation} Even to the quadratic order, the Fisher matrix is only an approximation to the curvature matrix of the likelihood function near the extrema. To examine the real shape of the likelihood function, we explicitly fit the curvature $\\sigma_B$ and log-normal offset $x_B$ and compared the results with the Fisher matrix (now diagonal) and the offsets suggested by \\citet{bond2000} (see their eq. [28]). In general we found very good agreement.  In a few bands, however, the uncertainty indicated by the true likelihood function is $\\sim 15\\%$ larger than that derived from the Fisher matrix. The uncertainty and log-normal offsets reported in this paper are determined from fits to the likelihood function.  \\subsubsection{Window Functions}\\label{subsubsec:window} Our goal is to produce a CMB power spectrum that can be used to constrain cosmological models. Window functions are used to convert a given model power spectrum to quantities directly comparable to the band-power measurements of the experiment. A window function $W_{B\\ell}$ for band ``$B$'' should have the following property: $$ \\langle q_B \\rangle=\\sum_\\ell (W_{B\\ell}/\\ell)D_\\ell\\,, $$ where $q_B$ is the experimental band-power measurement. \\citet{knox99} derived the appropriate window function for the band-power quadratic estimator \\citep{tegmark97b}, $$ W_{B\\ell}/\\ell=\\frac{1}{2}\\sum_{B'} (F^{-1})_{BB'}{\\rm Tr}\\left( \\frac{\\partial C_T}{\\partial D_\\ell}C^{-1}C_{T,B'}C^{-1}\\right)\\,. $$ This form of window function has been used in several other experiments \\citep{halverson_thesis,pryke02,myers02}, and should serve as a good approximation for maximum likelihood band-power. Making use of eqs.~[\\ref{fish}][\\ref{ctb}], it can be shown that the pre-decorrelation window functions satisfy the following normalization condition, \\begin{equation} \\sum_\\ell \\chi_{B\\ell}W_{B'\\ell}/\\ell=\\delta_{BB'}\\,.\\label{norm} \\end{equation} In Section~\\S\\ref{subsubsec:decorr}, we described the computation of a set of decorrelated band-powers from the Fisher matrix of the raw band-powers. This same linear transformation is applied to the window functions to produce a set of decorrelated window functions.  It is not necessary to calculate the window functions for all $\\ell$. However, we would like to sub-divide the $q_B$ bands into bins with finer $\\ell$ space resolution than the $q_B$ themselves. Otherwise in order to satisfy the normalization condition eq.~[\\ref{norm}], the pre-decorrelation window function reduces to a Kronecker-$\\delta$ in band index ``$B$'', or the tophat function given by $\\chi_{B\\ell}$. This is a reasonable approximation for experiments like BOOMERANG where the sky coverage, and therefore intrinsic $\\ell$ resolution, is large enough to resolve all the structures in the power spectrum. For {\\sc Acbar}, the $\\ell$ resolution of the experiment is comparable to the expected structure in $D_\\ell$ and it is essential to precisely characterize the dependence of each band power on the details of the power spectrum. The decorrelated window functions shown in Figure~\\ref{fig:window} were calculated with a resolution in $\\ell$ space of $\\Delta \\ell =30$. The first window function has significant oscillations with $\\Delta \\ell \\sim 110$. This appears to be an result of the close spacing of the fields in the LMT subtraction. For sufficiently smooth theoretical models these oscillations will average to a mean value, however, the exact dependence of the first band power on arbitrary input models could be complex. Numerical tabulations of the window functions are available on the {\\sc Acbar} public web site.   \\begin{figure*}[t] \\centerline{ \\psfig{figure=window_joint_pq_c3_14.ps,width=5in,height=3in,angle=0}} \\caption{The window functions ($W_{B\\ell}/\\ell$) for the decorrelated {\\sc Acbar} band powers. The vertical lines show the band boundaries. Numerical tabulations of these functions are given on the {\\sc Acbar} website. } \\label{fig:window} \\end{figure*}   \\subsubsection{Monte Carlo Simulations}  In order to test our analysis technique, we performed 300 Monte Carlo realizations of CMB2 and CMB5 maps that are consistent with our observed noise statistics and the same underlying $\\Lambda$CDM model. These maps were processed exactly as the real data, including using the same projection of corrupted modes, in order to test the ability of our pipeline to accurately determine the input CMB power spectrum. A quadratic estimator technique \\citep{tegmark97b} is used to rapidly determine the power spectrum for each of 300 realizations. In Figure~\\ref{fig:montecarlo}, we show the resulting average power spectrum overlaid on the input model. It is clear that the input model is recovered without bias, even in the low $\\ell$ bins that are sensitive to the angular scales on which the mode removal is occurring.   \\begin{figure*}[t] \\centerline{ \\psfig{figure=cl_joint_pq_c3_montecarlo.ps,width=6in,height=3in,angle=0}} \\caption{The results of 300 Monte Carlo runs using the measured {\\sc Acbar} noise correlation. The solid line is the input fiducial $\\Lambda$CDM model. It is clear that the analysis method accurately reproduces the input power spectrum.} \\label{fig:montecarlo} \\end{figure*} ",
        "conclusions": "\\label{sec:conclusion}  We have used the {\\sc Acbar} receiver to measure the angular power spectrum of the CMB at a frequency of $150\\,$GHz over the multipole range $\\ell = 100-3000$. The power spectrum we present is derived from approximately 21 weeks of Austral winter observations with {\\sc Acbar} installed on the $2.1\\,$m Viper telescope at the South Pole.  In the course of analyzing this data, we have employed new analysis techniques designed specifically for high sensitivity ground based CMB observations. Monte Carlo simulations are used to verify that the analysis method accurately recovers the power spectrum without bias. The power spectrum we present is robust and has passed stringent tests for systematic errors. Galactic dust emission and radio point sources do not contribute significantly to the observed power and are projected out in the final analysis. Although dusty protogalaxies cannot be ruled out as a source of confusion, the expected contribution to the measured power is negligible. Overall, the resulting power spectrum appears to be consistent with the damped acoustic oscillations expected in standard cosmological models. In a companion paper \\citep{goldstein02}, the {\\sc Acbar} power spectrum is used to place constraints on cosmological models. The power in the highest $\\ell$ {\\sc Acbar} bin is consistent with the excess power measured in the Deep CBI pointings. At this point, the {\\sc Acbar} data lack the sensitivity to place significant constraints on the origin of the excess power observed by CBI.  We acknowledge assistance in the design and construction of {\\sc Acbar} by the UC Berkeley machine and electronics shop staff. The support of Center for Astrophysics Research in Antarctica (CARA) polar operations has been essential in the installation and operation of the telescope. Percy Gomez, Kathy Romer, and Kim Coble are thanked for their assistance in monitoring the observations and telescope pointing. We thank Nils Halverson, Julian Borrill, and Radek Stompor for a careful reading of the draft and useful comments on analysis algorithms. Finally, we would like to thank John Carlstrom, the director of CARA, for his early and continued support of the project. The ACBAR program has been primarily supported by NSF office of polar programs grants OPP-8920223 and OPP-0091840. This research used resources of the National Energy Research Scientific Computing Center, which is supported by the Office of Science of the U.S. Department of Energy under Contract No. DE-AC03-76SF00098. Chao-Lin Kuo acknowledges support from a Dr. and Mrs. CY Soong fellowship and Marcus Runyan acknowledges support from a NASA Graduate Student Researchers Program fellowship. Chris Cantalupo, Matthew Newcomb and Jeff Peterson acknowledge partial financial support from NASA LTSA grant NAG5-7926.  \\appendix"
    },
    "0212/astro-ph0212130_arXiv.txt": {
        "abstract": "{  A new, momentum preserving fast Poisson solver for N-body systems sharing effective $O(N)$ computational complexity, has been recently developed by Dehnen (2000, 2002). We have implemented the proposed algorithms in a Fortran-90 code, and parallelized it by a domain decomposition using the MPI routines. The code has been applied to intensive numerical investigations of galaxy mergers, in particular focusing on the possible origin of some of the observed scaling relations of elliptical galaxies.  We found that the Fundamental Plane is preserved by an equal mass merging hierarchy, while it is {\\it not} in a scenario where galaxies grow by accretion of smaller stellar systems. In addition, both the Faber-Jackson and Kormendy relations are {\\it not} reproduced by our simulations. ",
        "introduction": "In the last few years, fast algorithms for computing N-body interactions relying on the general multipole expansion techniques have been developed.  These schemes, usually referred to as Fast Multipole Methods (FMMs, see, e.g., Greengard \\& Rokhlin 1987, 1997), are designed to have $O(N)$ computational complexity. However, it has been noted that only $O(N{\\rm log}N)$ scaling is usually achieved in numerical implementation (see, e.g., Capuzzo-Dolcetta \\& Miocchi 1998, hereafter CM98). The $O(N{\\rm log}N)$ operation count also characterizes the N-body codes based on tree algorithms, originally proposed by Barnes \\& Hut (1986, hereafter BH86). In addition, it has been shown by direct numerical tests that, for given accuracy, tree-codes are faster than the FMMs by a factor of $\\sim 4$ (CM98). For this reason, most of the N-body codes currently used for astrophysical applications are based on the classical BH86 tree scheme, where the particle data are organized into a nested oct-tree cell structure. Interactions between distant particles are approximated by cell--particle interactions, and the distribution of the particles on each cell is represented by a multipole expansion, usually truncated to the quadrupole term.  Recently, a new scheme has been introduced (Dehnen 2000, 2002, hereafter D02), which can be seen as an original combination of the tree-based BH86 scheme and of the FMMs multipole expansion, truncated at a fixed low order level (the proposed $p=3$ order results the optimal one). This scheme, implemented by the Author in a C++ code [named falcON, Force Algorithm with Complexity $O(N)$], results in a significant improvement over the existing BH86 tree-codes, in terms of uniform accuracy, linear momentum conservation, and CPU time performances. Moreover, the numerical tests presented in the D02 paper show that for an accuracy level $\\epsilon \\simeq 10^{-3}$, for the first time an {\\it effective} $O(N)$ scaling in operation count, as predicted by the FMM theory, is in fact obtained.  For these reasons two of us (P.L. and C.N.) have implemented this scheme in a new Fortran-90 code (FVFPS, a Fortran Version of a Fast Poisson Solver), and have then parallelized it by using the MPI procedures. A completion of the code to handle gas-dynamics using Euler equations and the AMR (Adaptive Mesh Refinement) techniques is under development.  Here we briefly describe the main characteristics of Dehnen's algorithm, and we present our FVFPS implementation and its parallelization. The achieved accuracy level, the performances in terms of operation count speed-up, the effective $O(N)$ scaling, as well as the parallel scaling with the processors number, are then documented in numerical tests. Finally, we describe some interesting astrophysical applications of the code, in the context of the study of galaxy merging.  Our version of the code is available upon request. ",
        "conclusions": "The main results of our work can be summarized as follows:  \\begin{itemize}  \\item Our implementation of the D02 scheme in a F-90 N-body code (FVFPS) has effective $O(N)$ complexity for number of particles $N \\gsim 10^4$, though its time efficiency is lower by a factor of $\\sim 1.5$ than Dehnen's C++ implementation.  \\item As shown by numerical tests, the parallelized version of the FVFPS code has good scaling with the number of processors (for example for numbers of particles $N \\simeq 10^6$, at least up to 16 processors).  \\item We found very good agreement between the results of merging simulations performed with our FVFPS code and with the GADGET code.  \\item From an astrophysical point of view, our high resolution simulations of hierarchies of {\\it dissipationless} galaxy mergers indicate that {\\it equal mass} merging is compatible with the existence of the FP relation. On the other hand, when an {\\it accretion} scenario is considered, the merger remnants deviate significantly from the observed edge--on FP (see also NLC02).  \\item The different behavior with respect to the FP of the merger remnants in the two scenarios is a consequence of structural non--homology, as shown by the analysis of their projected stellar density profiles.  \\item In any case, the end--products of our simulations fail to reproduce both the FJ and Kormendy relations. In the case of equal mass merging, the combination of large effective radii and low central velocity dispersions maintains the remnants near the edge--on FP.  \\end{itemize}"
    },
    "0212/astro-ph0212306_arXiv.txt": {
        "abstract": "Magnetic elements on the quiet sun are buffeted by convective flows that cause lateral motions on timescales of minutes. The magnetic elements can be observed as bright points (BPs) in the G band at 4305 {\\AA}. We present observations of BPs based on a long sequence of G-band images recorded with the Dutch Open Telescope (DOT) and post-processed using speckle masking techniques. From these images we measured the proper motions of isolated BPs and derived the auto-correlation function of their velocity relative to the solar granulation pattern. The accuracy of BP position measurements is estimated to be less than 23 km on the Sun. The rms velocity of BPs (corrected for measurement errors) is about 0.89 km s$^{-1}$, and the correlation time of BP motions is about 60 s.  This rms velocity is about 3 times the velocity measured using cork tracking, almost certainly due to the fact that isolated BPs move more rapidly than clusters of BPs.  We also searched for evidence of vorticity in the motions of G band BPs. ",
        "introduction": "Observations of the Sun with high spatial resolution show network bright points (Muller 1983, 1985, 1994; Muller \\& Keil 1983; Muller \\& Roudier 1984, 1992) and ``filigree'' (Dunn \\& Zirker 1973; Mehltretter 1974; Berger et al 1995), which are chains of bright features located within the intergranular lanes. The bright points and filigree are seen in the wings of strong spectral lines such as H$\\alpha$ and Ca II H \\& K, in lines formed in the photosphere, and even at continuum wavelengths (with reduced contrast). The widths of these structures is 100 to 200 km, at the limit of resolution of ground-based solar telescopes. In the following we collectively refer to these bright structures as bright points (BPs). The BPs are associated with regions of strong magnetic field (Chapman \\& Sheeley 1968; Title, Tarbell \\& Topka 1987; Simon et al 1988; Title et al 1989, 1992; Keller 1992) and correspond to magnetic flux tubes of kilogauss field strength that stand nearly vertically in the solar atmosphere (Stenflo 1973; Stenflo \\& Harvey 1985; S\\'{a}nchez Almeida \\& Mart\\'{i}nez Pillet 1994; see review by Solanki 1993).  The dynamical behavior of BPs has been studied by a number of authors. Muller (1983) found that facular points on the quiet sun are predominantly located in patches at the periphery of supergranule cells, indicating that the magnetic elements are advected by the supergranular flow. The BPs always first appear in the dark spaces at the junction of several granules, never inside a granule nor in the space between only two granules. As the granulation pattern evolves, the BPs remain in the intergranular spaces throughout their lifetime, but not necessarily at the junction of several granules like at the time of their first appearance. New BPs have a tendency to appear adjacent to existing points, and 15 \\% of the BPs seem to split into two points which move apart until a separation of 1 to {1.5\\arcsec} is reached.  Muller et al (1994) measured velocities of 29 isolated BPs and found a mean speed of 1.33 km s$^{-1}$. Strous (1994) studied BPs in a growing active region. Using line-center images taken in Fe~I (5576~{\\AA}), he found velocities between 0.26 km s$^{-1}$ and 0.62 km s$^{-1}$. Berger \\& Title (1996) measured velocities of 1 to 5 km s$^{-1}$ for G-band (4305~{\\AA}) BPs in the ``moat'' around a sunspot; they showed that the motions are constrained to the intergranular lanes and are primarily driven by the evolution of the granulation pattern. They found that the BPs continually split up and merge, with a mean time between splitting events of few hundred seconds. L\\\"{o}fdahl et al (1998) and Berger et al (1998) analyzed G-band and continuum images obtained at the Swedish Vacuum Solar Telescope (SVST) on La Palma and found rapid splitting and merging of BPs in an enhanced network region. Van Ballegooijen et al (1998) used a ``cork tracking'' method to measure the proper motions of G-band BPs and found that BPs appear to be passively advected by the granulation flow. The filigree are known to be associated with abnormal granulation patterns (Dunn \\& Zirker 1973), and the granules near network BPs are smaller and more numerous than near a normal intergranular space (Muller, Hulot \\& Roudier 1989), suggesting that the magnetic field has some effect on the granulation flow.  Little is known about the small-scale dynamics of flows in intergranular lanes and the interaction of magnetic elements with such flows. Theoretical models (e.g.,~Stein \\& Nordlund 2000; Emonet \\& Cattaneo 2001) predict that vorticity is concentrated in the lanes and that magnetic elements exhibit rapid rotational motions. At present, there is no direct observational evidence for vorticity in intergranular lanes. More detailed measurements of vorticity are needed for testing magneto-convection models, and as input for models of flux-tube waves that heat the upper solar atmosphere (Hasan et al 2002).  In this paper we present measurements of BP proper motions within the intergranular lanes. In particular, we look for evidence of vortical motions. Tracking BPs is a difficult task. Long sequences of images with very good seeing (or corrections of seeing) are needed. Image jitter due to residual atmospheric effects needs to be reduced as much as possible. The measurements of BP positions need to be made with respect to some frame of reference.  The solar granulation is used as a reference, and the motion of this reference frame is determined by correlation tracking on the granulation pattern. We also tested measuring the BPs with respect to the average position of the bright points in each field, but found that the result is nearly identical to the result from correlation tracking.  BPs vary widely in lifetime and contrast, so it is very difficult to automate the process of tracking them. In this paper, we use manual selection of BPs followed by fitting the peaks. ",
        "conclusions": "We have shown that with good enough data we can precisely measure positions for isolated BPs, obtaining both their linear and angular motions. The errors due to image jitter have been reduced to levels of 25 km (0.5~pixels) or less. This allows us to measure the motions of individual BPs to a precision of less than 1 km s$^{-1}$ for each 30 second time interval, which will be important for studying the generation of MHD waves in flux tubes. We also track BPs for sufficiently long times that we can see the regularity in their motions.  The rms velocities we measured for isolated BPs were almost a factor of three higher than velocities measured using cork tracking (van Ballegooijen et al 1998). We believe this is due to a selection effect. By only measuring isolated BPs we measure the fastest-moving features. Once BPs cluster with other BPs, mostly at lane interstices, they seem to be ``captured'' by the group and their motions are reduced. There are many more BPs in these clusters than there are isolated BPs so the cluster statistics dominate when measurements are made with cork tracking and the average velocity is reduced.  In measuring the motions of BPs, we have found that they move in somewhat circular paths. Combining the angular change of their motions with the distance they travel is a potential way of estimating the (vertical) vorticity in the flow field, assuming the BPs act as test particles. While performing the tracking of individual BPs, we saw little evidence of vorticity other than the tracks of the BPs themselves.  We looked for, but did not find, pairs of BPs that orbited one another; in fact, BPs that came in close proximity to other BPs tended to have reduced motion. We also saw little evidence of swirling motions in the granules or granular lanes that would correspond to the observed BP motions. This might indicate that BPs are affected not only by surface flows associated with the solar granulation, but also by other flows that occur at larger depth below the photosphere.  Further modeling will be required to establish whether or not the BPs move passively with the surface flow, and to determine whether the observed BP motions are consistent with models of flux tube waves and heating of the upper solar atmosphere."
    },
    "0212/astro-ph0212183_arXiv.txt": {
        "abstract": "We discuss in detail the pulsation properties of variable stars in globular clusters (GCs) and in Local Group (LG) dwarf galaxies. Data available in the literature strongly support the evidence that we still lack a complete census of variable stars in these stellar systems. This selection bias is even more severe for small-amplitude variables such as Oscillating Blue Stragglers (OBSs) and new exotic groups of variable stars located in crowded cluster regions. The same outcome applies to large-amplitude, long-period variables as well as to RR Lyrae and Anomalous Cepheids in dwarf galaxies. ",
        "introduction": "Variable stars in stellar systems such as GCs and dwarf galaxies have played a fundamental role in improving our knowledge on stellar populations (Baade 1958) as well as on the physical mechanism that drive the pulsation instability (Schwarzschild 1942). The main advantage of cluster variables when compared with field ones is that they are located at the same distance, and possibly the same reddening. Moreover, they formed from the same proto-globular cloud and therefore they have the same age, and chemical composition. Even though cluster variables present several undoubted advantages current knowledge concerning the pulsation properties of these objects is still limited. Recent estimates based on new data reduction procedures to perform differential photometry (ISIS, Alard 2000) suggest that the incompleteness factor in the detection of RR Lyrae stars is at least of the order of 30\\% (Kaluzny et al. 2001; Corwin \\& Carney 2001) in Galactic GCs characterized by high central densities. This limit is even more severe for OBSs, since the luminosity amplitude range from hundredths of a magnitude to a few tenths. Moreover, their radial distribution peaks toward the center of the cluster, and therefore ground based observations are strongly limited by crowding (Gilliland et al. 1998; Santolamazza et al. 2001). The same outcome applies to Miras and to Semi-Regular variables in GGCs, but for a different reason, quite often they are saturated in current CCD chips. This is a real limit for metal-rich clusters of the Galactic bulge, since they lack of RR Lyrae stars or host a few of them (Pritzl et al. 2002), and the detection of Miras could supply an independent distance estimate (Feast et al. 2002).  Variable stars in dwarf spheroidal (dSph) galaxies presents several pros and cons when compared with variables in GGCs. The star formation history as well as the dynamical evolution of dSph galaxies is much more complex than for GGCs. Typically the age of stellar populations in LG dSphs ranges from a few Gyr to 12-13 Gyr, i.e. as old as stars in GGCs (Da Costa 1999). Wide photometric surveys strongly support the evidence of extra-tidal stars near several dSphs (Irwin \\& Hatzidimitriou 1995; Martinez-Delgado et al. 2001). The observation of these stellar debris resembles the tidal tails detected in several GGCs (Leon et al. 2000). On the other hand, dSph galaxies apparently host large amounts of Dark Matter (DM), and indeed the mass-to-light ratios in these systems range from $(M/L)_V\\sim 5$ (Fornax) to $\\sim100$ (Ursa Minor). However, the scenario is still quite controversial and the evidence that dSphs present large DM central densities would suggest that they are not a large version of GGCs, since the latter present M/L ratios $\\approx1-2$. Photometric and spectroscopic data on variable stars in dSphs might supply new insights on the impact that environmental effects have on their evolutionary and pulsation properties. Unfortunately, data available in the literature are limited, since these stellar systems cover wide sky regions. The use of wide field imagers and wide field, multifiber spectrographs might overcome these problems.  In the following we discuss the impact that variables in stellar systems might have on cosmic distances and on stellar populations. ",
        "conclusions": "The results presented in the previous sections bring forward the evidence that the empirical scenario for variable stars in stellar systems such as GCs and LG dwarf galaxies is far from being complete. The limit applies not only to aperiodic variables but also to long-period variables such as Miras and Semi-Regulars. The same outcome applies to RR Lyrae stars affected by the Blazhko effect.  Even though several LG dSphs present stellar populations with chemical compositions and stellar ages quite similar to stars in GCs, the RR Lyrae variables present pulsation properties that are a {\\em bridge} between Oosterhoff type I and Oosterhoff type II clusters. This preliminary evidence seems to suggest that either the dynamical history and/or the chemical evolution in these stellar systems might play a role to explain this difference. In this context the use of the new wide field imagers will supply the unique opportunity to investigate on a star-by-star basis the stars and the variables in LG dwarf galaxies.  Although new theoretical frameworks have been developed to account for the occurrence of mixed-mode pulsators among RR Lyrae and OBSs we still lack a comprehensive explanation of the physical mechanisms that drive the occurrence of such a phenomena. It goes without saying that new sets of full amplitude nonlinear, convective models tightly connected with evolutionary models are highly requested to constrain the region of the instability strip where these pulsators present this intriguing behavior."
    },
    "0212/astro-ph0212460_arXiv.txt": {
        "abstract": "Experience with core-collapse supernova simulations shows that accurate accounting of total particle number and 4-momentum can be a challenge for computational radiative transfer. This accurate accounting would be facilitated by the use of particle number and 4-momentum transport equations that allow transparent conversion between volume and surface integrals in both configuration and momentum space. Such conservative formulations of general relativistic kinetic theory in multiple spatial dimensions are presented in this paper, and their relevance to core-collapse supernova simulations is described. ",
        "introduction": "\\label{sec:introduction}  The state of the art in core-collapse supernova simulations now includes energy- and angle-dependent neutrino transport \\cite{rampp00,mezzacappa01,liebendoerfer01b,rampp02,liebendoerfer02,thompson02,janka02,janka02b}. However, experience shows that simultaneous conservation of both energy and lepton number is difficult numerically \\cite{mezzacappa93b,liebendoerfer02}. This challenge motivates us to develop new conservative formulations of relativistic kinetic theory that are specifically attuned to the need for accurate accounting of particle number and energy in numerical simulations of radiative transfer problems.  Before describing the conservative formulations of kinetic theory we seek, we explain why energy- and angle-dependent neutrino transport is necessary in supernova simulations and detail the magnitude of the challenge of energy conservation.  Sophisticated treatments of neutrino transport are necessary because the ultimate energy source of the supernova explosion---the gravitational potential energy of the stellar progenitor's core---is eventually converted almost completely into neutrinos. Some of the gravitational potential energy is lost to escaping neutrinos during core collapse, but most of it is converted into a thermal bath of dense nuclear matter, photons, electron/positron pairs, and trapped neutrinos deep inside the nascent neutron star.  Neutrinos, having the weakest interactions, are the most efficient means of cooling; they diffuse outward on a time scale of seconds towards a semi-transparent region near the surface of the newly forming neutron star, and eventually escape with about 99\\% of the released gravitational energy. In modeling the conversion of gravitational potential energy into neutrino fluxes, energy- and angle-dependent neutrino transport is necessary to accurately follow the transition from quasi-isotropic diffusion to forward-peaked free streaming. In this transition region energy is transferred from the neutrino radiation to the matter behind a stalled shock wave, and this energy transfer may be necessary to propel the shock through the outer layers in an explosion \\cite{colgate66,bethe85}. But whether or not such neutrino heating is the proximate cause of explosion, the fact that neutrinos dominate the energetics implies that accurate neutrino transport is integral to any realistic and comprehensive study of the explosion mechanism.  The importance of accurate neutrino transport is a lesson learned from experience with supernova simulations. Parametrized studies \\cite{janka96} highlight the sensitivity of explosions to neutrino luminosities and to conditions in the semi-transparent region near the nascent neutron star's surface (see also Ref. \\cite{janka01})---precisely the region where neutrino energy and angle dependence must be tracked carefully. Moreover, there remains a nagging qualititative uncertainty in simulations with multidimensional hydrodynamics: Those with neutrino transport that depends on direction in configuration space but is averaged over energy and angles exhibit explosions \\cite{herant94,burrows95,fryer02}, while those with neutrino transport that depends on energy but is averaged over angles in both configuration and momentum space do not show explosions \\cite{mezzacappa98,mezzacappa98b}. It may be that these differing outcomes are due to the different neutrino transport schemes, both of which are ultimately inadequate.  Moving beyond these general arguments for the necessity of accurate neutrino transport, quantitative consideration of the energetics shows how severe the requirements are on one aspect of accuracy: energy conservation. As mentioned above, virtually all of the gravitational potential energy (\\inlineSFinmath${\\Tilde {{10}^{53}}\\multsp\\ \\Mvariable{\\rm erg}}$) released during collapse is eventually converted into intense neutrino fluxes lasting several seconds. However, supernova explosion energies (the kinetic energy of the ejecta) are observed to be only \\inlineSFinmath${\\Tilde {{10}^{51}}\\Mvariable\\ {\\rm erg}}$. Now, because it is difficult to argue with any rigor about the physical impact of any energy lost or gained due to numerical error, the total energy should be conserved to a precision corresponding to the phenomena of interest in the problem. Hence a simulation's result for explosion energy accurate to, say, \\inlineSFinmath${10\\InvisibleSpace \\%}$ would require total energy conservation to an accuracy of about \\inlineSFinmath${0.1\\InvisibleSpace \\%}$. Allowing for systematic error accrual, the total energy would have to be conserved at a level of \\inlineSFinmath${0.1\\InvisibleSpace \\%/N}$ per time step, where \\inlineSFinmath${N\\Tilde {{10}^5}}$ is the total number of time steps in the simulation.   Conservative formulations of kinetic theory would help to meet the numerical challenge of particle number and energy conservation in core-collapse supernova simulations. To give an idea of the kind of formulation of kinetic theory that we seek, we first review familiar descriptions of the dynamics of a fluid medium. The dynamics can be described in two different ways, which might be called {\\em elemental} and {\\em conservative}.  The elemental formulation expresses the evolution of the fluid in terms of equations of motion for the velocity and some independent set of quantities (e.g. temperature, densities of various species comprising the fluid, etc.) measured by an observer moving along with the fluid (``comoving observer''). For example, consider a spacetime with metric components $\\{g_{\\mu\\nu}\\}$ and metric determinant $g\\equiv \\det(g_{\\mu\\nu})$, containing a perfect fluid with  4-velocity components $\\{u^\\mu\\}$ and comoving frame total energy density $\\rho$, pressure $p$, and baryon density $n$. Iin the absence of radiative transfer and significant energy input from nuclear reactions, the perfect fluid evolves  according to \\begin{eqnarray} (\\rho +p){u^{\\mu }}\\Big(\\frac{\\partial {u^{i }}}{\\partial {x^{\\mu }}}+{{{{\\Gamma }^{i }}\\InvisibleSpace }_{\\Mvariable{\\rho \\mu }}}{u^{\\rho }}\\Big)+({g^{\\Mvariable{i \\mu }}}+{u^{i }}{u^{\\mu }})\\frac{\\partial p}{\\partial {x^{\\mu }}}&=&0, \\label{euler} \\\\ \\dispSFNumberedEquationmath{u^{\\mu }}\\frac{\\partial \\rho }{\\partial {x^{\\mu }}}+\\frac{(\\rho +p)}{{\\sqrt{-g}}}\\frac{\\partial }{\\partial {x^{\\mu }}}\\big({\\sqrt{-g}}{u^{\\mu }}\\big)&=&0, \\label{energy} \\\\ \\dispSFNumberedEquationmath{u^{\\mu }}\\frac{\\partial n}{\\partial {x^{\\mu }}}+\\frac{n}{{\\sqrt{-g}}}\\frac{\\partial }{\\partial {x^{\\mu }}}\\big({\\sqrt{-g}}{u^{\\mu }}\\big)&=&0.\\label{baryon} \\end{eqnarray} (Greek and latin letters are spacetime and space indices respectively.) Supplementary relations between \\inlineSFinmath${\\rho }$, \\inlineSFinmath${p}$, and \\inlineSFinmath${n}$---referred to as the {\\em equation of state}---are determined by the microphysics of the fluid. The name {\\em elemental} denotes the fact that by writing down separate equations of motion for the velocity and comoving-frame quantities, the kinetic and ``intrinsic'' fluid energies---two ``elements'' of the system---are analytically separated.  The conservative approach expresses the evolution of the system in terms of the divergence of the stress-energy tensor \\inlineSFinmath${{T^{\\Mvariable{\\mu \\nu }}}}$. For a perfect fluid, $T^{\\mu\\nu}\\inlineSFinmath{=(\\rho +p){u^{\\mu }}{u^{\\nu }}+p {g^{\\Mvariable{\\mu \\nu }}}}$, and Eqs. (\\ref{euler}) and (\\ref{energy}) are replaced by \\begin{equation} \\dispSFNumberedEquationmath{\\frac{1}{{\\sqrt{-g}}}\\frac{\\partial }{\\partial {x^{\\mu }}}\\big({\\sqrt{-g}}{T^{\\Mvariable{\\nu \\mu }}}\\big)=-{{{{\\Gamma }^{\\nu }}\\InvisibleSpace }_{\\Mvariable{\\rho \\mu }}}{T^{\\Mvariable{\\rho \\mu}}},}\\label{divergence} \\end{equation} while Eq. (\\ref{baryon}) is replaced by \\begin{equation} \\dispSFNumberedEquationmath{\\frac{1}{{\\sqrt{-g}}}\\frac{\\partial }{\\partial {x^{\\mu }}}\\big({\\sqrt{-g}}n\\multsp {u^{\\mu }}\\big)=0.}\\label{baryon2} \\end{equation} Volume integrals of the left-hand sides of Eqs. (\\ref{divergence}) and (\\ref{baryon2})---obtained by multiplying by the invariant spacetime volume element $\\sqrt{-g}\\,d^4x$ and integrating---are related to surface integrals in an obvious manner. Physically, this relates the time rates of change of 4-momentum and baryon number in a volume to fluxes through a surface surrounding that volume; hence the labeling of Eqs. (\\ref{divergence}) and (\\ref{baryon2}) as {\\em conservative}.  The right-hand side of Eq. (\\ref{divergence}) deserves note. After discussing the % reasons for this term's existence, we comment on what it means for conservation issues.  There are several important cases where the connection coefficients ${{{{\\Gamma }^{\\nu }}\\InvisibleSpace }_{\\Mvariable{\\rho \\mu }}}$ on the right-hand side of Eq. (\\ref{divergence}) might not vanish. They are present in curved spacetime, where (at least in part) they embody gravitational forces. But even in flat spacetime, coordinates employed by accelerated observers give rise to connection coefficients. And even without spacetime curvature or accelerated reference frames, connection coefficients arise from the use of curvilinear coordinates.  What does the right-hand side of Eq. (\\ref{divergence}) mean for 4-momentum conservation? Only when it vanishes---that is, only for inertial observers in flat spacetime employing rectilinear coordinates---are the components of total 4-momentum constant in time (``conserved''). For only in this case do the coordinates reflect the translation invariance of flat spacetime, the physical origin of 4-momentum conservation. (Curved spacetime lacks translation invariance, so there is no 4-momentum conservation.) Because the presence, for whatever reason, of a source term like the right-hand side of Eq. (\\ref{divergence}) means that the 4-momentum components in such a basis are not conserved, it might more properly be called a ``balance equation'' than a ``conservation law''. But because the volume integral of the left-hand side easily translates into a surface integral, we still call it a ``conservative'' formulation.  There are some special cases in which a conserved quantity associated with the time coordinate $t$ can be found, however. For unaccelerated observers in flat spacetime, ${{{{\\Gamma }^{t }}\\InvisibleSpace }_{\\Mvariable{\\rho \\mu }}}=0$, even in curvilinear coordinates. This means that energy is conserved, though the 3-momentum components in cuvilinear coordinates are not. Another special case is the restriction to spherical symmetry in general relativity: Here certain coordinate choices allow the non-vanishing ${{{{\\Gamma }^{t }}\\InvisibleSpace }_{\\Mvariable{\\rho \\mu }}}$ terms to be absorbed into the left-hand aside, leading to the identification of a conserved energy-like quantity (see e.g. Ref. \\cite{liebendoerfer01}).   Having used the familiar example of a perfect fluid to discuss what we mean by ``elemental'' and ``conservative'' formulations, we now consider kinetic theory in terms of these categories. The evolution of a particle type described by kinetic theory is often expressed as an equation of motion for the distribution function \\inlineSFinmath${f}$, the ensemble-averaged density of a given particle type in phase space. (These concepts will be defined with greater precision in subsequent sections; for the present discussion it is sufficient to understand that phase space is the combination of configuration space and momentum space, and that multiplying $f$ by the volume of an infinitesimal cell in phase space gives the number of particles having positions and momenta within the ranges defined by that cell.) The distribution function evolves due to advection through phase space and particle interactions.  Advection through phase space gives rise to derivatives of $f$ with respect to the components of the position vector $x$ and the momentum vector $p$. For numerical evolution it is convenient to parametrize the distribution function in terms of $\\{x^\\mu\\}$, the components of $x$ in a global ``coordinate basis'' \\footnote{In a ``coordinate basis'' the basis vectors are $\\vec{e}_\\mu = {\\partial\\over \\partial x^\\mu}$. These basis vectors ``commute'' because the order of partial derivatives can be interchanged freely. The connection coefficients in a coordinate basis are given by ${{{{\\Gamma }^{\\mu }}\\InvisibleSpace }_{\\Mvariable{\\nu \\rho }}}=\\frac{1}{2}{g^{\\Mvariable{\\mu \\sigma }}}\\Big(\\frac{\\partial {g_{\\Mvariable{\\sigma \\nu }}}}{\\partial {x^{\\rho }}}+\\frac{\\partial {g_{\\Mvariable{\\sigma \\rho }}}}{\\partial {x^{\\nu }}}-\\frac{\\partial {g_{\\Mvariable{\\nu \\rho }}}}{\\partial {x^{\\sigma }}}\\Big).$ In a ``non-coordinate basis'' obtained from the coordinate basis by a transformation $\\vec{e}_{\\mu'} = {L^\\mu}_{\\mu'}\\vec{e}_\\mu$, the basis vectors $\\vec{e}_{\\mu'} = {L^\\mu}_{\\mu'}{\\partial\\over \\partial x^\\mu}$ do not commute if ${L^\\mu}_{\\mu'}$ depends on position in spacetime. The connection coefficients in a non-coordinate basis have extra terms associated with the non-vanishing commutators of the basis vectors. In this paper, the connection coefficients in a non-coordinate basis are obtained from those of a coordinate basis via the transformation ${{{{\\Gamma }^{{\\mu' }}}}_{{\\nu' }{\\rho' }}}={{{{L }^{{\\mu' }}}}_{\\mu }}{{{{L }^{\\nu }}}_{{\\nu' }}}{{{{L }^{\\rho }}}_{{\\rho' }}}{{{{\\Gamma }^{\\mu }}}_{\\Mvariable{\\nu \\rho }}}+{{{{L }^{{\\mu' }}}}_{\\mu }}{{{{L }^{\\rho }}}_{{\\rho' }}}\\frac{\\partial {{{{L }^{\\mu }}}_{{\\nu' }}}}{\\partial {x^{\\rho }}}$. }. (Here and throughout this paper, quantities defined with respect to the coordinate basis have % indices without accents.) For momentum we make a different choice of basis, because interactions with a fluid are most easily described---and best handled numerically---in terms of momentum components measured by a comoving observer. The change from the coordinate basis to an orthonormal basis associated with the comoving observer (a ``non-coordinate basis'') has two parts. First there is a transformation \\inlineSFinmath${{{\\Mvariable{dx}}^{\\overvar{\\mu }{\\_}}}={{{e^{\\overvar{\\mu }{\\_}}}\\InvisibleSpace }_{\\mu }}{{\\Mvariable{dx}}^{\\mu }}}$ to an (in general non-comoving) orthonormal basis. (Here and throughout this paper, quantities defined with respect to the non-comoving orthonormal basis have indices accented with a bar.) This is followed by a Lorentz transformation \\inlineSFinmath${{{\\Mvariable{dx}}^{\\overvar{\\mu }{\\RawWedge }}}={{{{\\Lambda }^{\\overvar{\\mu }{\\RawWedge }}}\\InvisibleSpace }_{\\overvar{\\mu }{\\_}}}{{\\Mvariable{dx}}^{\\overvar{\\mu }{\\_}}}}$ to a comoving orthonormal basis. (Here and throughout this paper, quantities defined with respect to the comoving orthonormal basis have indices accented with a hat.) Hence advection through phase space will involve derivatives of $f$ with respect to the coordinate basis position components $\\{x^\\mu\\}$ and the comoving orthonormal basis momentum components $\\{p^{\\overvar{\\mu}{\\RawWedge}}\\}$.  Turning from advection to particle interactions, we consider the case where the particle species are sufficiently dilute that the interactions can be described in terms of a collision integral \\inlineSFinmath${\\mathbb{C}[f]}$ depending only on % the distribution functions $f$ of the individual particle species. (This approximation ignores correlations between particles; that is, the number of instances of finding particles at the same position is obtained from the product of their distribution functions.) In this case, the equation of motion for the distribution function \\inlineSFinmath${f}$ is \\cite{lindquist66,mezzacappa89} \\begin{equation} \\dispSFNumberedEquationmath{{p^{\\overvar{\\mu }{\\RawWedge }}}\\bigg({{{{\\Lambda }^{\\overvar{\\mu }{\\_}}}\\InvisibleSpace }_{\\overvar{\\mu }{\\RawWedge }}}{{{e^{\\mu }}\\InvisibleSpace }_{\\overvar{\\mu }{\\_}}}\\frac{\\partial f}{\\partial {x^{\\mu }}}-{{{{\\Gamma }^{\\overvar{\\nu }{\\RawWedge }}}\\InvisibleSpace }_{\\overvar{\\rho }{\\RawWedge }\\overvar{\\mu }{\\RawWedge }}}{p^{\\overvar{\\rho }{\\RawWedge }}}\\frac{\\partial f}{\\partial {p^{\\overvar{\\nu }{\\RawWedge }}}}\\bigg)=\\mathbb{C}[f].} \\label{boltzmann} \\end{equation} This is the {\\em Boltzmann equation}.   In terms of the categories described above in connection with fluid evolution, the Boltzmann equation is ``elemental'': The most fundamental quantity---the distribution function---is the evolved variable, and volume integrals of the equation in both configuration space and momentum space are not obviously related to surface integrals. This ``non-conservative'' character is present even in flat spacetime and rectangular coordinates. The first term of Eq. (\\ref{boltzmann}) is non-conservative even in flat spacetime and rectilinear coordinates because the boost \\inlineSFinmath${{{{{\\Lambda }^{\\overvar{\\mu }{\\_}}}\\InvisibleSpace }_{\\overvar{\\mu }{\\RawWedge }}}}$---which depends on the coordinates $\\{x^\\mu\\}$---sits outside the derivative $\\partial f/\\partial x^\\mu$. The spatial dependence of the boost also gives rise to non-vanishing ${{{{\\Gamma }^{\\overvar{\\nu }{\\RawWedge }}}\\InvisibleSpace }_{\\overvar{\\rho }{\\RawWedge }\\overvar{\\mu }{\\RawWedge }}}$, even in flat spacetime and rectilinear coordinates; therefore the second term of Eq. (\\ref{boltzmann}) does not vanish. This second term is non-conservative because the factor $p^{\\overvar{\\rho }{\\RawWedge }}$ sits outside the derivative ${\\partial f}/{\\partial {p^{\\overvar{\\nu }{\\RawWedge }}}}$.   While the Boltzmann equation is ``elemental'' or ``non-conservative'', it is well known (e.g., see Refs. \\cite{lindquist66,ehlers71,israel72}) the first two momentum moments of the distribution function (integrated over a suitable invariant momentum space volume element \\inlineSFinmath${\\Mvariable{dP}}$) constitute a particle number current \\inlineSFinmath${{N^{\\mu }}}$ and particle stress-energy tensor \\inlineSFinmath${{T^{\\Mvariable{\\mu \\nu }}}}$, \\begin{eqnarray} \\dispSFNumberedEquationmath{N^{\\mu }}&=&\\int f\\multsp {p^{\\mu }}\\Mvariable{dP} =\\int f\\multsp \\,{p^{\\overvar{\\mu }{\\RawWedge }}}{{{{\\Lambda }^{\\overvar{\\mu }{\\_}}}\\InvisibleSpace }_{\\overvar{\\mu }{\\RawWedge }}}{{{e^{\\mu }}\\InvisibleSpace }_{\\overvar{\\mu }{\\_}}}\\, \\Mvariable{dP}, \\label{particleNumberCurrent} \\\\ \\dispSFNumberedEquationmath{T^{\\Mvariable{\\mu \\nu }}}&=&\\int f\\multsp {p^{\\mu }}{p^{\\nu }}\\Mvariable{dP} = \\int f\\multsp \\, {p^{\\overvar{\\mu }{\\RawWedge }}}{{{{\\Lambda }^{\\overvar{\\mu }{\\_}}}\\InvisibleSpace }_{\\overvar{\\mu }{\\RawWedge }}}{{{e^{\\mu }}\\InvisibleSpace }_{\\overvar{\\mu }{\\_}}}\\, {p^{\\overvar{\\nu }{\\RawWedge }}}{{{{\\Lambda }^{\\overvar{\\nu }{\\_}}}\\InvisibleSpace }_{\\overvar{\\nu }{\\RawWedge }}}{{{e^{\\nu }}\\InvisibleSpace }_{\\overvar{\\nu }{\\_}}}\\, \\Mvariable{dP}, \\label{particleStressEnergy} \\end{eqnarray} which (for electrically neutral particles) obey the balance equations \\begin{eqnarray} \\dispSFNumberedEquationmath\\frac{1}{{\\sqrt{-g}}}\\frac{\\partial }{\\partial {x^{\\mu }}}\\big({\\sqrt{-g}}{N^{\\mu }}\\big)&=&\\int \\mathbb{C}[f]\\Mvariable{dP}, \\label{particleNumberConservation} \\\\ \\dispSFNumberedEquationmath\\frac{1}{{\\sqrt{-g}}}\\frac{\\partial }{\\partial {x^{\\mu }}}\\big({\\sqrt{-g}}{T^{\\Mvariable{\\nu \\mu }}}\\big) &=& -{{{{\\Gamma }^{\\nu }}\\InvisibleSpace }_{\\Mvariable{\\rho \\mu }}}{T^{\\Mvariable{\\rho \\mu }}} + \\int \\mathbb{C}[f]{p^{\\nu }}\\Mvariable{dP}. \\label{particleEnergyConservation} \\end{eqnarray}  While this result is often stated in the literature (e.g., see Refs. \\cite{lindquist66,ehlers71,israel72}), because of the non-conservative character of Eq. (\\ref{boltzmann}) it is not obvious how its momentum moments give rise to Eqs. (\\ref{particleNumberConservation}) and (\\ref{particleEnergyConservation}). Equation (\\ref{boltzmann}) contains factors ${{{{\\Lambda }^{\\overvar{\\mu }{\\_}}}\\InvisibleSpace }_{\\overvar{\\mu }{\\RawWedge }}}{{{e^{\\mu }}\\InvisibleSpace }_{\\overvar{\\mu }{\\_}}}$ outside ${\\partial f}/{\\partial {x^{\\mu }}}$; but from Eqs. (\\ref{particleNumberCurrent})-(\\ref{particleEnergyConservation}) we see that these factors have come inside the derivative with respect to $x^\\mu$. What happens to the spacetime derivatives of ${{{{\\Lambda }^{\\overvar{\\mu }{\\_}}}\\InvisibleSpace }_{\\overvar{\\mu }{\\RawWedge }}}{{{e^{\\mu }}\\InvisibleSpace }_{\\overvar{\\mu }{\\_}}}$ that are generated in taking these factors inside the derivative? Furthermore, according to Eqs. (\\ref{particleNumberCurrent}) and (\\ref{particleStressEnergy}), the quantities \\inlineSFinmath${{N^{\\mu }}}$ and \\inlineSFinmath${{T^{\\Mvariable{\\mu \\nu }}}}$ involve no momentum derivatives of $f$. But the second term of Eq. (\\ref{boltzmann}) contains a factor $p^{\\overvar{\\rho }{\\RawWedge }}({\\partial f}/{\\partial {p^{\\overvar{\\nu }{\\RawWedge }}}})$. Because of the momentum factor outside the momentum derivative it is not obvious how this term will contribute to Eqs. (\\ref{particleNumberConservation}) and (\\ref{particleEnergyConservation}) when integrated over momentum space.  How, then, is the connection between Eq. (\\ref{boltzmann}) and Eqs. (\\ref{particleNumberConservation}) and (\\ref{particleEnergyConservation}) established in detail? The reviews of relativistic kinetic theory by Lindquist \\cite{lindquist66} and Israel \\cite{israel72} do not provide detailed proofs. As discussed in Sec. \\ref{sec:distribution}, an explicit proof by Ehlers \\cite{ehlers71} relies on the fact that \\inlineSFinmath${{N^{\\mu }}}$ and \\inlineSFinmath${{T^{\\Mvariable{\\mu \\nu }}}}$ are momentum-integrated quantities, and no direct insight is gained into what happens to the momentum derivatives of $f$ in the integration over momentum space.     For those interested in computer models of radiative transfer problems, these are not idle academic questions; they are issues that must be faced in order to build simulations capable of making meaningful scientific statements about the core-collapse supernova explosion mechanism. Experience with supernova simulations in spherical symmetry shows that understanding the detailed connection between the Boltzmann equation and the particle number and 4-momentum balance equations has important consequences for how well these quantities are conserved in a simulation \\cite{mezzacappa93,liebendoerfer02}. While a discretization of the Boltzmann equation is a natural numerical method of evolving the neutrino species, na\\\"{\\i}ve differencings of the various terms in Eq. (\\ref{boltzmann}) generally will not be consistent with a straightfoward differencing of Eqs. (\\ref{particleNumberConservation}) and (\\ref{particleEnergyConservation}), leading to unacceptable numerical errors in particle number and energy conservation. % In Lagrangian (or ``comoving'') coordinates in spherical symmetry, Mezzacappa and Bruenn \\cite{mezzacappa93} derive a conservative formulation of the Boltzmann equation transparently related to particle number balance as expressed in Eq. (\\ref{particleNumberConservation}). They also devise methods of handling momentum derivatives that are consistent with both number and energy conservation \\cite{mezzacappa93b}. Liebend\\\"orfer et al. \\cite{liebendoerfer02} went a step further in this spherically symmetric case, deducing the connnection between the Mezzacappa and Bruenn ``number conservative'' Boltzmann equation in comoving coordinates and energy conservation as represented by an Eulerian (or ``lab frame'') version of Eq. (\\ref{particleEnergyConservation}). Using complicated, non-intuitive differencings of hydrodynamic and gravitational variables, they construct a numerical implementation of % the Mezzacappa and Bruenn ``number conservative'' Boltzmann equation that is stable and faithful to its analytic connection to the lab frame version of Eq. (\\ref{particleEnergyConservation}) to the accuracy necessary to make solid scientific statements about the neutrino-driven explosion mechanism in spherical symmetry \\cite{mezzacappa01,liebendoerfer01b,liebendoerfer02}.   Because spherically symmetric models with Boltzmann transport fail to reproduce some important observable characteristics of core-collapse supernovae (e.g. the launch of an explosion \\cite{mezzacappa01,liebendoerfer01b,liebendoerfer02,rampp00,rampp02,thompson02}), this work must be followed up in multiple spatial dimensions (see Refs. \\cite{janka02,janka02b} for some early efforts). In this paper we develop---allowing for full relativity and multiple spatial dimensions---conservative formulations of kinetic theory, such that volume integrals in both configuration and momentum space are {\\em trivially related} to surface integrals. These conservative expressions {\\em make transparent} the connection between Eq. (\\ref{boltzmann}) and Eqs. (\\ref{particleNumberConservation}) and (\\ref{particleEnergyConservation}). They can be used to deduce the term-by-term cancellations involved in this connection, thereby illuminating the complicated differencings required to achieve the cancellations numerically. Our conservative formulations also suggest new primary variables of radiation transport: particle number and energy variables describing the contribution of each {\\em comoving orthonormal basis} momentum bin to the {\\em coordinate basis} number and energy densities. It may be that the use of these new radiation variables could provide a simpler path to an accurate accounting of particle number and energy in simulations of radiative transfer problems.  The organization of this paper is as follows. Differential forms and exterior calculus are natural mathematical tools for handling the volume elements and integrations needed in relativistic kinetic theory. We closely follow (and slightly extend) Ehlers' work \\cite{ehlers71} in reviewing their use in the description of phase space for particles of definite mass (Sec. \\ref{sec:phaseSpace}) and the derivation of the Boltzmann equation (Sec. \\ref{sec:distribution}). The centerpieces of this paper---two conservative reformulations of the Boltzmann equation, which provide transparent connections to particle number and energy-momentum balance as expressed in Eqs. (\\ref{particleNumberConservation}) and (\\ref{particleEnergyConservation})---are presented in Sec. \\ref{sec:conservative}. Because differential forms and exterior calculus may be unfamiliar to those whose primary interests are radiation transport, in general, or supernova science, in particular, Secs. \\ref{sec:phaseSpace}-\\ref{sec:conservative} each will contain two subsections: one containing a general derivation, and a second that explicitly demonstrates aspects of the derivation in the familiar case of the $O(v)$ limit in Lagrangian coordinates in spherical symmetry. (While we review some aspects of exterior calculus in our presentation, these are mostly in endnotes, and are more in the character of reminders than a self-contained introduction. For the latter we refer the reader to Refs. \\cite{ehlers71,mtw}.) A conclusion (Sec. \\ref{sec:conclusion}) summarizes our results, comments on their connection to moment formalisms, and discusses the utility of these formulations for supernova simulations. As an application of our formalism, an appendix contains \\inlineSFinmath{$O (v)$} equations in flat spacetime, but in coordinates sufficiently general to represent rectangular, spherical, and cylindrical coordinate systems. ",
        "conclusions": "\\label{sec:conclusion}  In this section we summarize our conservative formulations of kinetic theory, comment on their relation to moment formalisms, and discuss their possible application in the core-collapse supernova environment.  Having in mind computational radiative transfer in astrophysical environments, we have sought formulations of relativistic kinetic theory with the following properties: (1) they are expressed in terms of global, Eulerian (or ``lab-frame'') spacetime coordinates \\inlineSFinmath${\\{{x^{\\mu }}\\}}$; (2) they are expressed in terms of convenient three-momentum coordinates \\inlineSFinmath${\\big\\{{u^{\\overvar{i}{\\RawWedge }}}\\big\\}}$ (e.g. spherical polar), which are taken from the orthonormal momentum components \\inlineSFinmath${\\big\\{{p^{\\overvar{i}{\\RawWedge }}}\\big\\}}$ measured by an observer comoving with the medium; and (3) they are ``conservative'', having transparent connections to total particle number and 4-momentum balance as expressed in Eqs. (\\ref{numberConservation}) and (\\ref{energyConservation}).  To express our formulations having these properties, we here introduce the {\\itshape specific particle number flux vector } \\begin{equation} \\dispSFNumberedEquationmath{{{\\ScriptCapitalN }^{\\mu }}\\equiv {{{{\\ScriptCapitalL }^{\\mu }}\\InvisibleSpace }_{\\overvar{\\mu }{\\RawWedge }}}{p^{\\overvar{\\mu }{\\RawWedge }}}f} \\end{equation} and the {\\itshape specific particle stress-energy tensor} \\begin{equation} \\dispSFNumberedEquationmath{{{\\ScriptCapitalT }^{\\Mvariable{\\mu \\nu }}}\\equiv {{{{\\ScriptCapitalL }^{\\mu }}\\InvisibleSpace }_{\\overvar{\\mu }{\\RawWedge }}}{{{{\\ScriptCapitalL }^{\\nu }}\\InvisibleSpace }_{\\overvar{\\nu }{\\RawWedge }}}{p^{\\overvar{\\mu }{\\RawWedge }}}{p^{\\overvar{\\nu }{\\RawWedge }}}f,} \\end{equation} where the transformation to the comoving frame \\inlineSFinmath${{{{{\\ScriptCapitalL }^{\\overvar{\\mu }{\\RawWedge }}}\\InvisibleSpace }_{\\mu }}}$ is given by equation (\\ref{compositeTransformation}). (While the adjective ``specific'' often denotes a quantity measured per unit mass, in this context we use it to denote the particle flux and stress-energy in a given invariant momentum space volume element.) While the distribution function \\inlineSFinmath${f}$ of a given particle type of mass \\inlineSFinmath${m}$ and charge \\inlineSFinmath${e}$ obeys the Boltzmann equation (Eq. (\\ref{fullBoltzmann})), the specific particle number flux and stress-energy satisfy the  conservative equations \\begin{eqnarray} \\frac{1}{{\\sqrt{-g}}}\\frac{\\partial }{\\partial {x^{\\mu }}}\\big({\\sqrt{-g}}{{\\ScriptCapitalN }^{\\mu }}\\multsp \\big)+\\nonumber \\\\ E({\\bf p})\\Bigg|\\det \\big[\\frac{\\partial {\\bf p}}{\\partial {\\bf u}}\\big]{{\\Bigg|}^{-1}}  \\frac{\\partial }{\\partial {u^{\\overvar{i}{\\RawWedge }}}}\\Bigg(\\frac{1}{{E({\\bf p})}}\\Bigg|\\det \\big[\\frac{\\partial {\\bf p}}{\\partial {\\bf u}}\\big]\\Bigg|\\big(e\\multsp {{{F^{\\overvar{j}{\\RawWedge }}}\\InvisibleSpace }_{\\overvar{\\mu }{\\RawWedge }}}-{{{{\\Gamma }^{\\overvar{j}{\\RawWedge }}}\\InvisibleSpace }_{\\overvar{\\mu }{\\RawWedge }\\overvar{\\nu }{\\RawWedge }}}{p^{\\overvar{\\nu }{\\RawWedge }}}\\big)\\frac{\\partial {u^{\\overvar{i}{\\RawWedge }}}}{\\partial {p^{\\overvar{j}{\\RawWedge }}}}{{{{\\ScriptCapitalL }^{\\overvar{\\mu }{\\RawWedge }}}\\InvisibleSpace }_{\\mu }}{{\\ScriptCapitalN }^{\\mu }}\\Bigg)\\IndentingNewLine \\nonumber\\\\ = \\mathbb{C}[f], \\label{conservative1} \\\\ \\frac{1}{{\\sqrt{-g}}}\\frac{\\partial }{\\partial {x^{\\nu }}}\\big({\\sqrt{-g}}{{\\ScriptCapitalT }^{\\Mvariable{\\mu \\nu }}}\\big)+ \\nonumber \\\\ {E({\\bf p})}\\Bigg|\\det \\big[\\frac{\\partial {\\bf p}}{\\partial {\\bf u}}\\big]{{\\Bigg|}^{-1}} \\frac{\\partial }{\\partial {u^{\\overvar{i}{\\RawWedge }}}}\\Bigg(\\frac{1}{{E({\\bf p})}}\\Bigg|\\det \\big[\\frac{\\partial {\\bf p}}{\\partial {\\bf u}}\\big]\\Bigg|\\big(e\\multsp {{{F^{\\overvar{j}{\\RawWedge }}}\\InvisibleSpace }_{\\overvar{\\nu }{\\RawWedge }}}-{{{{\\Gamma }^{\\overvar{j}{\\RawWedge }}}\\InvisibleSpace }_{\\overvar{\\nu }{\\RawWedge }\\overvar{\\rho }{\\RawWedge }}}{p^{\\overvar{\\rho }{\\RawWedge }}}\\big)\\frac{\\partial {u^{\\overvar{i}{\\RawWedge }}}}{\\partial {p^{\\overvar{j}{\\RawWedge }}}}{{{{\\ScriptCapitalL }^{\\overvar{\\nu }{\\RawWedge }}}\\InvisibleSpace }_{\\nu }}{{\\ScriptCapitalT }^{\\Mvariable{\\mu \\nu }}}\\Bigg)\\IndentingNewLine \\nonumber\\\\  = {{{F^{\\mu }}\\InvisibleSpace }_{\\nu }}{{\\ScriptCapitalN }^{\\nu }} - {{{{\\Gamma }^{\\mu }}\\InvisibleSpace }_{\\Mvariable{\\nu \\rho }}}{{\\ScriptCapitalT }^{\\Mvariable{\\nu \\rho }}} + {{{{\\ScriptCapitalL }^{\\mu }}\\InvisibleSpace }_{\\overvar{\\mu }{\\RawWedge }}}{p^{\\overvar{\\mu }{\\RawWedge }}}\\mathbb{C}[f].\\label{conservative2} \\end{eqnarray} In stating that Eqs. (\\ref{conservative1}) and (\\ref{conservative2}) constitute  conservative formulations of particle kinetics, we mean that the connection to the balance equations of Eqs. (\\ref{numberConservation}) and (\\ref{energyConservation}) is transparent, in the following sense. We can use elementary calculus to form the familiar invariant momentum space volume element \\begin{equation} \\dispSFNumberedEquationmath{\\frac{{d^3}p}{E({\\bf p})}=\\frac{1}{{E({\\bf p})}}\\Bigg|\\det \\big[\\frac{\\partial {\\bf p}}{\\partial {\\bf u}}\\big]\\Bigg|{d^3}u,}\\label{momentumElement2} \\end{equation} where a transformation from orthonormal momentum components \\inlineSFinmath${\\big\\{{p^{\\overvar{i}{\\RawWedge }}}\\big\\}}$ to some other set of coordinates (e.g. momentum space spherical coordinates \\inlineSFinmath${\\big\\{{u^{\\overvar{i}{\\RawWedge }}}\\big\\}=\\{|{\\bf p}|,\\vartheta ,\\varphi \\}}$) has been performed. Multiplying Eqs. (\\ref{conservative1}) and (\\ref{conservative2}) by Eq. (\\ref{momentumElement2}) and integrating, the terms with momentum space derivatives are obviously transformed into vanishing surface terms; the results are Eqs. (\\ref{numberConservation}) and (\\ref{energyConservation}) for total particle number and 4-momentum balance.    In terms of differential forms, the procedure for obtaining the  conservative formulations of kinetic theory is straightforward. First, express the volume element \\inlineSFinmath${\\Omega }$ in the phase space for particles of definite mass in terms of the desired spacetime and 3-momentum coordinates. Next, by contraction with the Liouville vector, form the hypersurface element \\inlineSFinmath${\\omega ={L_m}\\cdot \\Omega }$. Then bring the exterior derivative \\inlineSFinmath${d(f\\multsp \\omega )}$ into the form \\inlineSFinmath${\\mathbb{N} [f]\\Omega }$ by direct computation; Eq. (\\ref{conservative1}) results on comparison with the Boltzmann equation. This result can then be used in conjunction with an evaluation of \\inlineSFinmath${d({v_{\\mu }}{p^{\\mu }}f\\multsp \\omega )}$ for arbitrary \\inlineSFinmath${{v_{\\mu }}}$ to obtain Eq. (\\ref{conservative2}). The reason the procedure is straightforward is that the ``heavy lifting'' of transforming the Boltzmann equation into  conservative forms is handled by two ``levers'' of considerable power: the generalized Stokes theorem, and the key relation \\inlineSFinmath${\\Mvariable{d\\omega }=0}$, which is closely related to the relativistic Liouville theorem.  A concrete example of our formalism is provided in the appendix, which contains \\inlineSFinmath{$O (v)$} equations for the specific particle number density, specific particle energy density, and their angular moments---all in flat spacetime, but in coordinates sufficiently general to represent rectangular, spherical, and cylindrical coordinate systems.   We now comment on the connection of Eqs. (\\ref{conservative1}) and (\\ref{conservative2}) to moment formalisms. In the usual treatments, if one writes the distribution function as a function of momentum variables as measured in a given frame (lab or comoving), it is natural to form moments by multiplying the distribution function by, for example, energies and angles measured in that frame, and integrating. Lo and behold, it turns out that these moments are number densities and fluxes, and energy/momentum densities and fluxes: components of a particle number flux vector and stress-energy tensor, measured in the same frame chosen to measure angles and energies. Traditional, then, are treatments in which the components of conserved tensors as measured in a given frame are expressed as functions of momentum variables as measured in that same frame.  But this traditional approach to moments may not be the most convenient, and Eqs. (\\ref{conservative1}) and (\\ref{conservative2}) provide an attractive alternative. For example, Liebend\\\"orfer et al. \\cite{liebendoerfer01,liebendoerfer02} form moments in the traditional way, resulting in components of conserved tensors as measured in an {\\em orthonormal comoving frame.} But it is the tensor components in the {\\em lab frame} that one would like to check, either because it is the coordinate basis (at least in a spatially multidimensional simulation) and therefore natural to deal with, or because it is the basis in which energy is conserved (in a simulation in comoving coordinates in spherical symmetry). The required transformation of the tensor components between these frames leads to the numerical complexity mentioned in Sec. \\ref{sec:introduction}, and discussed further in the latter parts of Subsecs. \\ref{subsec:conservativeGeneral} and \\ref{subsec:conservativeComoving}. In contrast, when the coordinate basis is in the lab frame as expected in spatially multidimensional simulations, {\\itshape integration of Eqs. (\\ref{conservative1}) and (\\ref{conservative2}) over momentum variables leads directly to the tensor components in the desired coordinate basis, even though the specific particle number and stress-energy are functions of comoving frame momentum variables.} The insight here is that the frame in which the tensor components are measured need not be the same as that employed to obtain the momentum variables used to parametrize the particle distributions. The appendix provides an example of a moment formalism of this kind.   In simulations of systems like core-collapse supernovae---in which careful attention to energetics is critical---a number of possible approaches, based on the conservative formulations of kinetic theory presented in this paper, might be suggested.  First, the general approach of, for example, Refs. \\cite{mezzacappa01,liebendoerfer01b,liebendoerfer02} could be followed, in which the conservative particle number distribution equation (Eq. (\\ref{conservative1})) is solved, and the conservative particle energy equation (the time component of Eq. (\\ref{conservative2})) is used as a consistency check. The quantitiy $\\sqrt{-g}\\,{\\cal N}^0$, which might be called the {\\em specific particle number density}, would be the primary neutrino distribution variable: It is the contribution of each {\\em comoving frame} momentum bin to the {\\em lab frame} particle number density. This approach makes number conservation a somewhat natural outcome, but energy conservation would require finite-differenced representations of various quantities to be ``matched'' in order that Eq. (\\ref{matching}) be satisified numerically. This might be considered the most rigorous and self-consistent method. (Note that if one solves the ``plain'', non-conservative Boltzmann equation---Eq. (\\ref{fullBoltzmann})---for the scalar distribution function $f$ as a function of comoving frame momentum variables, conservation of {\\em neither} lab frame particle number or energy is straightforward. The same is true of methods (e.g., Ref. \\cite{burrows00}) based on a non-conservative form of the transport equation for the comoving frame specific intensity.)  In order to avoid the intricate finite differencing of numerous terms demanded by this method, a second option would be the use of \\inlineSFinmath$\\sqrt{-g}\\, {{{\\ScriptCapitalT }^{00}}}$, which might be called the {\\itshape specific particle energy density, }as the primary neutrino distribution variable. Designing a code around a differenced version of the \\inlineSFinmath${\\mu =0}$ component of Eq. (\\ref{conservative2}) would make accurate accounting of total neutrino energy (as represented by the \\inlineSFinmath${\\mu =0}$ component of Eq. (\\ref{energyConservation})) relatively straightfoward. Of course, the neutrino number balance equation expressed in terms of \\inlineSFinmath$\\sqrt{-g}\\,{{{\\ScriptCapitalT }^{00}}}$ would have numerical errors unless certain finite differencings were carefully designed. But with respect to the crucial energetics of the physical system, it is worth noting that there are a couple of factors mitigating the impact of errors in number conservation in comparison with errors in energy conservation. Errors in number conservation translate into errors in the electron fraction \\inlineSFinmath${{Y_e}}$, which affect energy conservation through the equation of state, but: (1) Only \\inlineSFinmath${{{\\nu }_e}}$ and \\inlineSFinmath${{{\\left( \\overvar{\\nu }{\\_} \\right) }_e}}$ affect \\inlineSFinmath${{Y_e}}$, while all species impact the energy budget. Better to have two species contributing to error rather than six! (2) The effects of \\inlineSFinmath${{{\\nu }_e}}$ and \\inlineSFinmath${{{\\left( \\overvar{\\nu }{\\_} \\right) }_e}}$ on \\inlineSFinmath${{Y_e}}$ are opposite in sign (unlike their contributions to energy), so that to the extent that their distributions are similar, the impact of their errors on \\inlineSFinmath${{Y_e}}$ may approximately cancel.    A third possibility would be to solve for {\\itshape both} the specific particle energy density \\inlineSFinmath$\\sqrt{-g}\\,{{{\\ScriptCapitalT }^{00}}}$ and the specific particle number density \\inlineSFinmath$\\sqrt{-g}\\,{{{\\ScriptCapitalN }^0}}$. (Rampp and Janka \\cite{rampp02} solve for both number and energy distributions, but in the comoving frame; this limits the utility of their approach with respect to accurate tracking of lab frame quantities.) Instead of pre-defining both the boundaries and center values of bins in energy space, one could define the boundaries only and use the values of $\\sqrt{-g}\\,{{{\\ScriptCapitalT }^{00}}}$ and $\\sqrt{-g}\\,{{\\ScriptCapitalN }^0}$ (along with the transformation \\inlineSFinmath${{{{{\\ScriptCapitalL }^{\\overvar{\\mu }{\\RawWedge }}}\\InvisibleSpace }_{\\mu }}}$) to obtain center values of the energy bins in each spatial zone and each time step. The consistency of the solutions would be arguably reasonable as long as the derived center values of the energy bins do not wander outside the pre-defined bin boundaries.            \\appendix*"
    },
    "0212/astro-ph0212526_arXiv.txt": {
        "abstract": "{We derive the galaxy luminosity functions in V-, R-, I-, and K-bands of the cluster EIS\\,0048 at $z \\sim 0.64$ from data taken at the ESO Very Large Telescope. The data span the restframe wavelength range from UV, which is sensitive to even low rates of star formation, to the NIR, which maps the bulk of the stellar mass.  By comparing our data and previous results with pure luminosity evolution models, we conclude that bright ($M \\leq M^*+1$) cluster galaxies are already assembled at $z \\sim 1$ and that star formation is almost completed at $z \\sim 1.5$. ",
        "introduction": "One of the main issues in theoretical and observational research is to understand the physical processes driving the formation and evolution of bright massive galaxies in clusters, and to constrain the relative time scales.  The evolution of the cluster galaxy sequence (the colour--magnitude relation) has been studied for large cluster samples up to redshift $z \\sim 1$ (Aragon--Salamanca et al. \\cite{AEC93}; Lubin \\cite{LUB96}; Ellis et al. \\cite{ESD97}; Stanford et al. \\cite{SED98}, hereafter SED98; Nelson et al. \\cite{NGZ01}, hereafter NGZ01).  The results are consistent with a monolithic collapse scenario (see e.g. Larson \\cite{L74}) in which galaxies form at high redshifts and subsequently undergo passive evolution.  However, the evolution of colours only inform on the epoch when the bulk of stellar mass formed, while cluster bright galaxies ($M \\leq M^*+1$, mostly early--type) could have been also assembled recently ($z < 1$) from mergers of smaller units, at least as long as the merging processes do not induce strong star formation.  A different approach consists in the study of the evolution of the cluster luminosity function (CLF). In particular, the near--infrared (NIR) CLF can be used to assess the assembly history of galaxies, because the NIR light mainly informs on galaxy mass.  From the analysis of the values of $M^*$ in the K-band as a function of redshift for a sample of 38 clusters in the range $0.1 < z < 1$, De Propris et al. (\\cite{DPSE99}, hereafter DPSE99) found a trend consistent with passive luminosity evolution (PLE) (see also Nakata et al. \\cite{NAK01}, hereafter NKY01).  DPSE99 conclude that the mass function of bright cluster galaxies is invariant at $z < 1$, and that the assembly of those galaxies is largely complete by $z \\sim 1$. Spectroscopic data also show that massive galaxies already exist at least up to $z=0.83$ (van Dokkum et al. \\cite{vDF98}).  In this work we derive the CLFs for EIS\\,0048 at $z \\sim 0.64$ in the V-, R-, I-, and K-band (UV to NIR restframe).  The cluster membership has been assessed with the photometric redshift technique up to $M-M^* \\sim 1.5-2.5$ (according to the different depth of each band).  Since no other selection has been applied, the samples are not biased toward a particular galaxy population.  The paper is organized as follows. In Sect. 2 we introduce the photometric data, discuss the background subtraction, the selection of cluster members, the completeness of the samples, and obtain the CLFs. In Sect. 3 we discuss the results in terms of galaxy formation and evolution. In Sect. 4 we give a summary of the paper. In this work we assume H$_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M=0.3$ and $\\Omega_{\\Lambda} = 0.7$. ",
        "conclusions": "\\label{CONC}  We derived the galaxy LF in V-, R-, I-, and K-bands of the cluster EIS\\,0048 at $z \\sim 0.64$ from new photometric observations carried out at ESO VLT with FORS2 and ISAAC. Cluster members have been selected with the photometric redshift technique in order to enhance the contrast among cluster- and background--foreground galaxies. The remaining field contamination has been estimated as detailed in LBMI03.  The CLFs have been obtained up to $M-M^* \\sim 1.5-2.5$ (according to the waveband). We modelled the CLFs with a weighted parametric fit of the Schechter function, fixing the faint end slope to be $\\alpha = -0.9$. The values of $M^*$ obtained in I- and K-bands are in agreement with the analysis of DPSE99 and NGZ01 at similar redshifts.  We collected results from literature and introduced PLE models with different formation redshifts in order to discuss the evolution of $M^*$ in the V- and K-bands.  The key factors driving the evolution of bright galaxies in the scenario of hierarchical structure formation are the epoch when the bulk of stellar populations is formed, the cosmological time when mergers are effective to assemble the galaxies, the amount of star formation induced by mergers, and the age of the youngest stars.  Since the evolution of $K^*$ is consistent with a PLE scenario at least up to $z \\sim 1$, we can conclude that mergers must have already assembled bright ($M \\leq M^*+1$) cluster galaxies at this redshift (see also DPSE99 for a thorough discussion).  At $z = 0.64$ the V-band samples the UV restframe wavelength region and is sensitive to even low levels of SF.  We find that the SF in bright cluster galaxies has to be almost completed at $z \\sim 1.5$, whereas the formation redshift is $z_f \\geq 3$ assuming $\\tau \\sim 1$ Gyr as the time scale for SF. These results lead to conclude that the structure and the stars of bright cluster galaxies must have been formed between $z=4\\pm 1$ and $z=1.2\\pm0.2$."
    },
    "0212/hep-th0212327_arXiv.txt": {
        "abstract": "We consider the quasi-de Sitter geometry of the inflationary universe. We calculate the energy flux of the slowly rolling background scalar field through the quasi-de Sitter apparent horizon and set it equal to the change of the entropy (1/4 of the area) multiplied by the temperature, $dE=T dS$. Remarkably, this thermodynamic law reproduces the Friedmann equation for the rolling scalar field. The flux of the slowly rolling field through the horizon of  the quasi-de Sitter geometry is similar to the accretion  of a rolling scalar field onto a black hole, which we also analyze.  Next we add inflaton fluctuations which generate scalar metric perturbations. Metric perturbations result in a variation of the area entropy. Again, the equation $dE=T dS$ with fluctuations reproduces the linearized Einstein equations. In this picture as long as the Einstein equations hold, holography  does not put limits on the quantum field theory during inflation. Due to the accumulating metric perturbations, the horizon area during inflation randomly wiggles with dispersion increasing with time. We discuss this in connection with the stochastic decsription of inflation. We also address the issue of the instability of inflaton fluctuations in the ``hot tin can'' picture of de Sitter horizon. ",
        "introduction": "\\label{sec:intr}  The inflationary paradigm established during the last 20 years assumes that the primordial equation of state is almost vacuum-like: $p\\approx -\\epsilon$.  To realize this equation of state, most models deal with a scalar field $\\phi(t)$  (or other fields which in combination act as an effective scalar field) slowly rolling to the minimum of its potential $V(\\phi)$. During the slow roll regime the homogeneous scalar field produces geometry which can be well approximated by the quasi-de~Sitter metric.  The full pure de~Sitter spacetime, which corresponds to a 4d hyperboloid of constant curvature,  can be compactly represented by its Penrose diagram, given by the full square in Fig. \\ref{fig:penrose}. It can be covered by different coordinates. Cosmologists most often use coordinates in which the metric is time-dependent  and corresponds to an expanding flat universe \\begin{equation}\\label{flat} ds^2= -dt^2 + e^{ 2Ht} \\left(dr^2+r^2 d\\Omega^2 \\right) \\ , \\end{equation} where $d\\Omega^2=d \\theta^2+\\sin^2 \\theta d \\phi^2$. This coordinate system covers the upper  half of the hyperboloid, which  corresponds to the expansion branch.  The Penrose diagram of de Sitter spacetime in flat FRW coordinates is shown on the left panel of Fig.~\\ref{fig:penrose}. Quasi-de~Sitter geometry is described by the scale factor $a(t)=a_0 \\, e^{\\int dt H(t)}$, where the Hubble parameter $H$ is a slowly varying function of time, $\\dot H \\ll H^2$.  \\begin{figure}[b] \\centerline{\\epsfig{file=ds2b.eps, height=5cm}\\hspace{2cm} \\epsfig{file=ds3.eps, height=4.5cm}} \\medskip \\caption{ Penrose diagram of de~Sitter spacetime in the flat FRW coordinates (left) and the static coordinates (right). Each point represent a sphere $S^2$. Its radius at the horizon (dashed line on the left, edge of diamond on the right) is equal to $\\frac{1}{H}$. } \\label{fig:penrose} \\end{figure}  The time-dependent form of the metric (\\ref{flat}) is very convenient for investigating the dynamics of a scalar field with the equation \\begin{equation}\\label{scalar} \\Box \\phi=V_{,\\phi} \\end{equation} and for quantizing this field in the de~Sitter spacetime \\cite{bd}. Among quantum scalar fields with mass $m$ and conformal coupling $\\xi$ in de~Sitter geometry, the case of minimal coupling $\\xi=0$ and very small mass $m \\ll H$ plays an especially important role. Indeed, the regularized vacuum expectation value is $\\langle \\delta \\phi^2\\rangle =\\frac{3H^4}{8\\pi^2 m^2}$.  Formally, as was noted before the discovery of inflation, this is an odd case since its eigen-spectrum contains an infrared divergent term: $\\langle \\delta \\phi^2\\rangle \\to \\infty$ as $m \\to 0$. On the other hand, this is the most interesting case for application to  inflation, since the theory of inflaton  (as well as tensor) fluctuations is reduced exactly to this case.  Following the time evolution of individual  fluctuations, it  was found  that the infrared divergence can be interpreted  as  the  instability of quantum fluctuations of  a very light  scalar field, which are accumulated with time $\\langle \\delta \\phi^2 \\rangle =\\frac{H^3}{4\\pi} \\, t$ \\cite{fluc,linde,star}. Fluctuations of $\\delta \\phi$ induce scalar metric perturbations \\cite{metr,star,hawking,guth,bardeen}. This picture is a basis of the inflationary paradigm so successfully confirmed observationally. Notice that heavy or conformal fields are not produced by inflation. Further, backreaction of fluctuations $\\delta \\phi$ leads to the picture of stochastic evolution of quasi-de~Sitter geometry \\cite{stoch1,stoch2}, and at large values of $H$ even to self-reproduction (eternal) of the inflationary universe \\cite{self}. Scalar field in the eternal inflationary universe is described naturally in terms of the probability distribution function $P(\\phi, t)$ \\cite{stoch2,stoch3}.   Recently, de~Sitter spacetime and inflation have drawn significant attention in the theoretical physics/superstring community. Some of the most interesting topics are holography and the thermodynamics associated with the de~Sitter horizon.  In this context, the static form of the metric of the de Sitter  spacetime \\begin{equation}\\label{static} ds^2=-(1-H^2 \\, R^2)d\\tau^2+ {{dR^2} \\over{(1-H^2 \\, R^2)}}+R^2 d\\Omega^2 \\ \\end{equation} is commonly used. The Penrose diagram of de Sitter spacetime in static coordinates is plotted on the right  panel of Fig.~\\ref{fig:penrose}. The classical result of \\cite{GH} is that observer at the origin detects a thermal radiation from the de Sitter horizon at $R=\\frac{1}{H}$ with the temperature $T=\\frac{H}{2\\pi}$, and the horizon area $A=\\frac{4\\pi}{H^2}$ is associated with the huge (geometrical) entropy $S=\\frac{A}{4G}$. Thermal vacuum in the causal patch (``hot tin can'') corresponds to the Bunch-Davies vacuum of the metric (\\ref{flat}) \\cite{BMS,kleban} and gives a complementary picture of scalar field(s) fluctuations. It is not clear to us, however, how quantum fluctuations in the ``hot tin can'' picture correspond to the instability of quantum inflaton fluctuations $\\delta \\phi$ and generation of metric perturbations. We will return to this point at the end of the paper.  One of the issues in the holographic approach is the bookkeeping of entropy of de~Sitter spacetime. The holography bound declares that the geometrical entropy of the horizon exceeds the entropy of quantum states (of fields and particles) within the volume surrounded by the horizon. It was recently claimed that counting the entropy of quantum fluctuations generated during inflation in the ``hot tin can'' and comparing it to the change of the apparent horizon entropy violates the holography bound  unless an ultra-violet cutoff of order of $\\sim 10^{16}$ GeV in the momenta of fluctuations is imposed \\cite{kaloper}.  While it is expected that the approaches based on the time-dependent form of the de Sitter metric with unstable fluctuations and the static form of the de Sitter metric with thermal flux should give us complementary insights, their languages are apparently different. This is partly due to the difference between quasi-de~Sitter and  pure de~Sitter geometries, and partly because different questions are addressed. However, we have to understand how these two different approaches to (quasi-)de~Sitter geometry with a scalar field   are compatible with each other with respect to such important issues as the generation of fluctuations,  entropy and global geometry.  In this paper we consider a particular question of how  the apparent horizon area $A$, or the entropy $S=\\frac{A}{4G}$, vary due to the slow roll of the background scalar field and the generation of scalar metric perturbations  during inflation. A novel element here is that we combine the concepts of a dynamical, slowly rolling background field and the instability of its fluctuations, with the concept of geometrical, holographic  entropy.  In Section \\ref{sec:geom}, we will calculate a variation of the geometrical  entropy  due to the energy flux through the apparent horizon area. We find that, remarkably, the thermodynamical relation $\\delta E=TdS$ is equivalent to the Einstein equation for the rolling inflaton field. In a sense, our  derivation of a  correspondence between thermodynamics and the Einstein equations for inflation is a realization of such a correspondence found in an inspiring paper \\cite{jacob} for local accelerating observers. However, we introduce a technique to treat the apparent horizon of $R \\times S^2$ topology which is different from the description \\cite{jacob} of a local Rindler horizon for an accelerating observer.  As we will see, a non-vanishing flux is generated by the  kinetic term $\\dot \\phi^2$ of the slowly rolling inflaton field.  It turns out that this problem is very similar to the problem of the interaction of a homogeneous rolling  scalar field  with a runaway potential $V(\\phi)$ and a black hole.  In Section \\ref{sec:bh}, we switch our attention from inflation to black holes. A rolling scalar field  interacting with a black hole is a  transparent illustration of the energy  flow of a light scalar field through a horizon.  In Section \\ref{sec:fluc1}, we return to inflation. On top of the rolling background inflaton,  we consider  inflaton  fluctuations $\\delta \\phi$, which generate scalar metric perturbations $\\Phi$. We study the energy flux through the horizon including inhomogeneous $\\delta \\phi$ fluctuations  and corresponding variations in the area of the horizon, or entropy $dS$, which are sensitive to the scalar metric perturbations $\\Phi$. In this case, the calculations are more involved than the calculations for the homogeneous time dependent background field in  Section \\ref{sec:geom}. This happens because there is no exact Killing vector generating the horizon. However, for metric perturbations which preserve  spherical symmetry we still can define $TdS$ and compare it with the energy flux through the horizon.  In Section \\ref{sec:fluc2}, we argue that the  metric perturbations generated from inflaton quantum fluctuations can indeed be treated as (locally) spherically symmetric. We apply the general formalism for spherically symmetric non-static geometry  of Section \\ref{sec:fluc1} to fluctuations from inflation. Again, we find that the thermodynamical relation $\\delta E=TdS$ leads to equations connecting $\\Phi$ and $\\delta \\phi$ which are in exact agreement with the linearized Einstein equations for the fluctuations from inflation. This allows us to give new insights into the entropy of cosmological fluctuations, as we will discuss in Section \\ref{sec:disc}. ",
        "conclusions": "\\label{sec:disc}  In this paper we calculated the apparent horizon area and geometrical entropy of quasi-de~Sitter geometry, which describes an inflationary stage, and the energy flux of the scalar fields  through the apparent horizon.  Issues related to de Sitter thermodynamics  are commonly considered in the  static de Sitter coordinates, where an observer at the origin is surrounded by the event horizon and detects a thermal flux of temperature $T=\\frac{H}{2\\pi}$.  We work with the geometrical entropy $S$  of the apparent horizon  in the time-dependent planar de-Sitter coordinates adopted in inflationary cosmology, which admit a simple generalization to the quasi-de~Sitter geometry and, most importantly, admit a clear interpretation of fluctuations generated from inflation.  The slow roll of the inflaton field leads to slow change in the horizon radius. Assuming that the energy flux of the rolling scalar field through the horizon  changes the geometrical entropy $\\delta E=TdS$, we reproduce the Einstein equation (\\ref{einst})  which relates $\\dot H$ and $\\dot \\phi^2$. Other background scalars $\\chi$ which are subdominant during inflation give similar contributions to the energy flux $\\dot \\chi^2$, independent of their potentials.  This change of the quasi-de~Sitter horizon radius due to the rolling scalar field is very similar to the increase of the mass of a black hole due to the accretion of a background scalar field rolling  towards the minimum of its potential $\\dot M \\sim \\dot \\phi^2$. This type of accretion, which can be described analytically with the delayed field approximation, can be realized for runaway potentials of the background scalar field or for very light massive fields. We present this material here mainly to illustrate similar calculations  for rolling scalars in quasi-de~Sitter geometry, as the astrophysical effects of accretion of a rolling scalar onto a black hole that we considered (say quintessence and astrophysical-size black holes) are negligibly small. Oscillating heavy scalar field accretes onto black hole very differently as quasi-particles.  Inflationary, quasi-de~Sitter stage erases classical inhomogeneities which could exist prior to it. However, quantum fluctuations of the inflaton field (as well as other light scalars minimally coupled to gravity) are unstable and produce long-wavelength fluctuations of the inflaton field $\\delta \\phi$ which behave as classical inhomogeneities at scales larger than the Hubble patch of size $H^{-1}$. Fluctuations $\\delta \\phi$ generate long-wavelength scalar  metric perturbations $\\Phi$. Geometrical quantities, like an area of the apparent horizon, acquire corrections due to the scalar metric perturbations. We calculate the energy flux of the inhomogeneous scalar  field $\\phi(t)+\\delta \\phi(t, \\vec x)$ through the apparent horizon and the change in the apparent horizon area of the perturbed metric multiplied by its (geometrical) temperature $T dS$. Again, equating $\\delta E=TdS$, we show that this thermodynamics relation is compatible with the linearized Einstein equations which relate $\\Phi$ and $\\delta \\phi$.  Thus, as long as the Einstein equations hold, generation of the inflaton fluctuations is in full agreement with the variation of the entropy of the quasi-de~Sitter horizon. Therefore, in the picture of rolling inflaton with quantum fluctuations generated with time, we did not find that quantum fluctuations may violate holographic bound  during inflation. (Notice that the fluctuations from inflation are described by squeezed states, which do not carry entropy \\cite{bpm,Kiefer}. In simple terms, locally the effect of fluctuations is just a  wiggling of the local Hubble parameter).    It was suggested in \\cite{kaloper} that in the ``hot tin can'' picture the entropy of fluctuations may violate holographic bound  during inflation.\\footnote{If frozen fluctuations from inflation were carrying large entropy, then    a number $N$ of free light scalars produced during inflation would have $N$ times bigger entropy and UV cutoff proposed in \\cite{kaloper} would depend on $N$.} The issue of inflaton fluctuations in the ``hot tin can'' picture is not clear to us. In theory of inflation, scalar field fluctuations are usually considered in the time dependent metric (\\ref{flat}), as it was described in Section \\ref{sec:fluc2}. Bunch-Davies vacuum corresponds to the thermal state in the static coordinates. On the other hand, in the static de~Sitter coordinates, discussion of the quantum field theory is usually restricted to thermal radiation from the horizon, usually in terms of the detector response. Thermal radiation associated with the de~Sitter horizon  is related to any free quantum field, scalar fields with mass $m$ and any coupling to gravity, vector fields etc., the difference will be only in the radial dependence of the wave functions (``gray-body factor''). On the other hand, instability of quantum fluctuations from inflation occurs only for very light minimally coupled scalar fields. It will be interesting to understand what is relation of the quantum field theory in  the ``hot tin can'' picture,  and the instability of fluctuations from inflation. For this it will be essential to work with regularized VEV $<\\delta \\phi^2>$.  As we already mentioned, in the picture of rolling scalar field with accumulating inflaton  fluctuations, which we adopt in the paper, holographic bound and quantum fluctuations from inflation are compatible.  A  lesson from our calculations is that in the quasi-de~Sitter geometry with slowly rolling scalar field, the horizon area, or geometrical entropy, are perturbed by the scalar metric fluctuations according to the formulas (\\ref{wiggles}), (\\ref{wiggles2}). This means that the area of the horizon for an ensemble of the different Hubble patches is not the same but is a statistical variable by itself. For small metric perturbations (of order of $\\Phi \\sim \\frac{H}{M_p}$), its mean value is $\\bar A=\\frac{4\\pi}{H^2}$, where $H(t)$ is slowly decreasing background value, but dispersion around the mean value is defined by $\\langle\\Phi^2\\rangle$, which  growth with time. Equivalently, we can talk about statistical properties of the local Hubble values $\\tilde H$  \\cite{star,stoch3}. Notably, this picture is converging with the stochastic description of inflation in terms of the probability distribution of the inflaton field \\cite{stoch1,stoch2,stoch3}. Probability to have the value of inflaton field $\\phi$ in quasi-static regime is $P(\\phi)\\sim \\exp{ \\left(-\\frac{3M_p^4}{8V(\\phi)}\\right)}$ and can be interpreted in terms of entropy $S=\\log P(\\phi)$ \\cite{entr}. Remarkably, this entropy is identical to the geometrical entropy  $\\frac{A}{4G}$. During inflation and rolling of $\\phi$ the Wiener process of accumulating fluctuations  $\\delta \\phi$ (or similarly perturbations of  $\\Phi$) changes  distribution of $P(\\phi)$. Distribution of the horizon areas of different Hubble patches, defined by the local values of the Hubble parameters, depends on the background slow roll regime and the regime of accumulation of fluctuations, both of which depend on the model of inflation \\cite{stoch1,stoch2,stoch3}. It would be interesting to understand further the correspondence between stochastic approach to inflation and geometrical entropy (\\ref{wiggles}), (\\ref{wiggles2}).    So far we discussed small metric fluctuations $\\Phi$ or small local variations of $H$. However, in the chaotic inflationary scenario for large enough values of $\\phi$, variations of $H$ due to the quantum jumps of $\\delta \\phi$ can be large and lead to the self-reproducing inflationary universe \\cite{self}.  We would like to draw attention to this regime (which is still below the Planck energy density) and to note that ``adiabatic'' geometrical thermodynamics which  we considered in this paper is not applicable here. Indeed,  consider the Hubble patch where $H$ is increasing due to the quantum jumps. Increase of $H$ is not compatible with the classical Einstein equation (\\ref{einst}). Consequently, quantum jumps are not compatible with the horizon thermodynamics (of a single Hubble patch) since $\\delta E=TdS$ does not hold either. Geometrical entropy of the local Hubble patch is decreasing. However, we have to take into account the entropy of all Hubble patches. Self-reproduction of inflating regions looks like a chain reaction, which is described by the branching diffusion process \\cite{branch}."
    },
    "0212/astro-ph0212018_arXiv.txt": {
        "abstract": "A new hybrid approach to air shower simulations is described. At highest energies, each particle is followed individually using the traditional Monte Carlo method; this initializes a system of cascade equations which are applicable for energies such that the shower is one-dimensional.  The cascade equations are solved numerically down to energies at which lateral spreading becomes significant, then their output serves as a source function for a 3-dimensional Monte Carlo simulation of the final stage of the shower.  This simulation procedure reproduces the natural fluctuations in the initial stages of the shower, gives accurate lateral distribution functions, and provides detailed information about all low energy particles on an event-by-event basis.  It is quite efficient in computation time. ",
        "introduction": "The field of highest energy cosmic rays is an exciting subject with many open questions: What is the nature of the primary cosmic ray? What are the highest energies? What are possible sources/acceleration mechanisms? Is there clustering of events? Is there a GZK cutoff due to the microwave background? Ongoing (HIRES, AGASA) and future (Auger, OWL, EUSO) cosmic ray experiments aim to shed light on these mysteries.  At these high energies, direct measurement of the primary cosmic ray is impossible due to the low flux, which is only of order one event per square-kilometer per century at the highest observed energies. But cosmic rays initiate showers in the atmosphere, a cascade of secondary particles from collisions with air molecules, which themselves collide and so on. Experiments measure these air showers and reconstruct from their properties information about the primary ray at the beginning of the reaction.  Air shower models are of crucial importance for the reconstruction of the energy and primary type. The straightforward approach is to model each possible interaction of hadrons, leptons and photons with air molecules, and trace all secondary particles. At high energies this leads quickly to unpractical computation times, since the time grows with the number of particles in the shower and therefore increases rapidly with the primary energy.  A shower of even \\( 10^{19}eV \\) has more than $10^{10}$ particles at its maximum and would take months to compute. The thinning algorithm proposed by Hillas \\cite{Hillas1985} tries to solves this problem: below a fraction \\( f_{\\mathrm{thin}} \\) of the primary energy only a small sample of the particles is actually followed in detail, attributing them a higher weight. This procedure introduces artificial fluctuations and one must compromise between these and computation time.  People have tried to overcome these difficulties by defining systems of (mostly one-dimensional) transport equations which describe air showers \\cite{Kalmykov:1986ii,GaisserBook}. The numerical solutions of these equations can then be combined with a Monte-Carlo in order to account for natural fluctuations due to the first interactions and for lateral spread of low-energy particles \\cite{Dedenko1968,Lagutin:1999xh}. This is the principle of the hybrid method. Another realization of the hybrid approach is to use shower libraries in which presimulated longitudinal profiles are combined to compute the one-dimensional properties of air showers \\cite{Gaisser1997,Alvarez-Muniz:2002ne}.  In a recent paper \\cite{Bossard:2000jh}, a new approach to an old idea was introduced: the method of cascade equations, which allows one to compute longitudinal characteristics of air showers numerically in a very short time.  In this paper we introduce the further development of this approach. Traditional Monte-Carlo methods are combined with cascade equations in a hybrid approach. This allows to construct an efficient model which accounts not only for natural fluctuations due to the first interactions but also for the correct 3-dimensional spreading of low-energy secondary particles. In reasonable computing time it is possible to calculate longitudinal profiles and lateral distribution functions with detailed knowledge about particle momenta and arrival-times, which are reliable on an event-by-event basis. ",
        "conclusions": ""
    },
    "0212/astro-ph0212532_arXiv.txt": {
        "abstract": "{ We present the results of a search for globular clusters in the surroundings of 15 low surface brightness dwarf galaxies belonging to the Fornax Cluster, which was carried out on CCD images obtained with the $C$ and $T_1$ filters of the Washington photometric system. The globular cluster candidates show an extended and probably bimodal $(C-T_1)$ color distribution, which is inconsistent with the presence of a single population of metal-poor clusters detected in several dwarf galaxies. The surface number density of these candidates shows no concentration towards the respective dwarf galaxies, in whose outskirts they have been identified. On the contrary, if we split the candidates in two groups according to their projected distances to the center of the Fornax Cluster, those located closer to the center show a higher projected density than those located farther from it. These results suggest that the potential globular clusters might not be bound to the dwarf galaxies. Alternatively, these globulars could form part of the very peripheral regions of \\object{NGC 1399} (the central galaxy of the Fornax Cluster) or even belong to the intracluster medium. ",
        "introduction": "The Fornax Cluster is an excellent target for studying globular clusters: it is very rich, contains different types of galaxies and the globular cluster candidates can be detected at least as unresolved sources. There are numerous photometric studies of globular cluster systems in selected Fornax galaxies, particularly the central one \\object{NGC 1399} as well as other early-type galaxies. Most of these studies show that the color distribution of the globular clusters is bimodal due to the presence of two globular cluster populations, ``red\" and ``blue\"; these integrated colors are mainly driven by metallicity in objects as old as these ones.  With regard to the globular cluster system of \\object{NGC 1399}, the Washington photometry by \\citet{ost93} confirmed the existence of a color gradient that had been detected earlier by \\citet{bri91}, and suggested that the color distribution was multimodal, as was also supported by the V,I photometry by \\citet{kis97}. The HST imaging by \\citet{for98} and \\citet{gri99}, and a refined analysis of their previous Washington photometry by \\citet{ost98}, pointed to the bimodal character of the color distribution. More recently, the wide-field study by \\citet{di02a,di02b} showed that globular custers with a broad metallicity distribution -- that cannot be fitted with a single Gaussian -- extend further than 100 kpc from \\object{NGC 1399} center.  There are fewer investigations of globular cluster systems in dwarf galaxies that are not in the Local Group. Miller and collaborators \\citep{mi98a,mi98b,mil99,lot01} carried out a survey with images from the Wide Field Planetary Camera 2 of the Hubble Space Telescope (FOV $\\approx$ 6~arcmin$^2$) to analyze the properties of globular clusters and nuclei of dwarf elliptical galaxies (dEs) in the Fornax and Virgo Clusters and the Leo Group. They include about 20 dEs from the Fornax Cluster but none of them are in common with the present work. They show that the globular cluster specific frequency $S_\\mathrm{N}$ (number of globular clusters with respect to the parent galaxy's luminosity) of the dEs is in the range 2--6; the luminosity function  of  the globular cluster candidates is  consistent with a  Gaussian with peak at $M_V^0 \\approx - 7$ mag and dispersion $\\sigma_V \\approx~1.4$ mag (assuming a distance modulus of 31.4 for the Fornax Cluster); and most of the globular cluster $(V-I)$ colors are similar to those of the metal-poor globular cluster population \\citep[also][]{ash93}. The globular cluster system of the luminous dE,N galaxy \\object{NGC 3115 DW1} was studied by \\citet{du96a} and \\citet{puz00} who obtained  mean metallicities $[Fe/H]= -1.2$ and $-1$, respectively, and they both agreed on a specific frequency $S_\\mathrm{N}$ = 4.9. \\citet{du96b} obtained, on the basis of Washington photometry, a low mean metallicity, $[Fe/H]= -1.45$, for the globular cluster systems of two dE galaxies in the Virgo Cluster and they suggested that the dwarf galaxies globular cluster systems present a similar range of metallicities ($[Fe/H] = -2$ to $-1$) as globular clusters in the halos of spiral galaxies, in agreement with \\citet{ash93}. Turning to the Local Group, \\citet{min96} constructed a master-dE galaxy combining the data from globular cluster systems of several dEs in this Group; they found an old and metal-poor globular cluster population whose metallicity distribution matched the one of the Milky Way halo globulars.  The abovementioned bimodality in the color distribution of globular clusters in elliptical galaxies, is closely related to the formation of the globular clusters and a variety of scenarios have been proposed (for reviews on the subject see, e.g.,  \\citealt{ash98} or \\citealt{har01}). \\citet{ash92} and \\citet{zep93} predicted this bimodal metallicity distribution of globulars in elliptical galaxies as a result of gas-rich  mergers; they proposed that the blue population originally belonged to the progenitor galaxies and the red population formed during the merger. The numerical simulations from \\citet{bek02} showed that dissipative mergers create new globular clusters but they were not able to reproduce properly the bimodal metallicity distribution observed in elliptical galaxies.  A different point of view was exposed, e.g.,  by \\citet{for97} who found a correlation between the mean metallicity of the metal-rich globular clusters and the parent galaxy luminosity, which suggested that they share the same chemical enrichment process, while the mean metallicity of the metal-poor ones seems to be independent of the galaxy luminosity. They proposed that the bimodality originated in two phases of globular cluster formation from gas with different metallicities, and that most of them formed \"in situ\". The HST study of 17 early-type galaxies by \\citet{lar01} showed a correlation between the colors of both, blue and red globular clusters populations, with the B-luminosity and central velocity dispersion of the host galaxy, and concluded that their observations support globular cluster formation \"in situ\", in the protogalaxy potential well.  Alternatively, \\citet{for01} analyzed the relation between the mean color of blue and red globulars with the galaxy velocity dispersion and suggested that red globular clusters share a common origin with the host galaxy and blue ones seem to have formed quite independently; according to \\citet{cot98} these blue globular clusters may have been captured from other galaxies by merger processes or tidal stripping.  The idea of the accretion of dwarf galaxies into the cD halo of \\object{NGC 1399}, the stripping of their gas and globular clusters and the formation of new clusters from this gas poses a different origin for the red globular clusters \\citep{hil99}. \\citet{bur01} analyzed the blue globular cluster populations from 47 galaxies and found no correlation between their mean metallicity, which is very similar for all these systems, and the galaxy properties (luminosity, velocity dispersion, etc); they proposed that the metal-poor globular clusters may have formed from gas fragments of similar metallicity, as already suggested by \\citet{ash93}, and located within the dark halo of the galaxy.  More recently, the semianalytic model by \\citet{bea02} assumed that the metal-poor globular clusters formed in protogalactic fragments and the metal-rich ones originated in the gas-rich mergers of such fragments that occurred later.  Assuming the presence of globular clusters inside clusters of galaxies, an alternative scenario is proposed by \\citet{whi87} and \\citet{wes95}, who pointed to the possible existence of a population of globular clusters that are not bound to individual galaxies; instead, they are supposed to move freely in the central regions of the galaxy clusters. These intracluster globular clusters may be the result of interactions or mergers between the galaxies, or they may have formed precisely in the environment of a galaxy cluster without any parent galaxy. The kinematic analysis by \\citet{min98} and by \\citet{kis99} also suggested that some globular clusters may be associated with the gravitational potential of the galaxy cluster and not solely with \\object{NGC 1399}. On the other hand, \\citet{gri99} found no evidence of intergalactic globular clusters in an HST/WFPC image at a radial distance of about $1\\,\\fdg4$ from \\object{NGC 1399}, but due to the small field of view, they were not able to rule out their existence. Several objections against the intraclusters were raised by \\citet{har98} who tried to explain by means of their existence the supposed high specific frequency of M87, the central Virgo galaxy; but the latest values of $S_\\mathrm{N}$ obtained for NGC 1399 by \\citet{ost98} and \\citet{di02b} showed that it is not so high ($S_\\mathrm{N}$ = 5.6 and 5.1, respectively).  In favor of the existence of intergalactic material, \\citet{the97}, \\citet{men97}, and \\citet{cia98} presented evidence for the presence of intergalactic planetary nebulae within the Fornax and Virgo Clusters, while \\citet{fer98} reported several hundred of intracluster red giants in Virgo. Some globular clusters may have been stripped with them from other cluster galaxies if this is their origin \\citep{har01}.  In this paper, we analyze the characteristics of globular cluster candidates found near dwarf galaxies in the Fornax Cluster and its connection with the above mentioned scenarios. It is organized as follows: Section 2 describes the observations and the adopted criteria for the globular cluster candidates' selection. In Sect. 3 we analyze the color distribution, luminosity function and spatial distribution of the candidates. Finally, a summary of the results and a discussion on their implications are provided in Sect. 4. Preliminary results of this work have been presented by \\citet{bas02}. ",
        "conclusions": "With regard to the color distribution of the globular cluster candidates, it is interesting to note that we do not find a single population, the metal-poor one, as seems to be the common case for dwarf galaxies \\citep{du96b,mi98a,mi98b,mil99},  but an extended distribution, which appears to be bimodal though we cannot prove it statistically due to the small sample involved. In addition, if we take into account the specific frequency estimated for dwarf galaxies \\citep{mi98a,elm99} and the luminosity function of the globular cluster candidates, we should have found significantly fewer globular cluster candidates than we actually do.  According to the projected density of the potential globular clusters, they show no concentration towards the dwarfs while they do show concentration with respect to the center of the cluster. These results led us to speculate that the globular cluster candidates may not be associated to the dwarf galaxies themselves. We are then left with three possibilities: first, that these globular cluster candidates belong to the globular cluster system of \\object{NGC 1399}; second, they may be moving freely throughout the potential well of the cluster, without being bound to any galaxy in particular; or third, that they are a mix of both. In the first case, we should be accepting that the globular cluster system of \\object{NGC 1399} is much more extended and numerous than previously thought. According to this hypothesis, Figure\\,\\ref{f6} suggests that there should be clusters up to at least an intermediate angular distance of 80\\arcmin~(a projected distance of about 430 kpc) from the central galaxy; the CCD study over the largest area around \\object{NGC 1399} was performed by \\citet{di02b}, which extends up to 22\\arcmin, and showed that the globular cluster system extends over a radial projected distance of more than 100 kpc.   The total number (background corrected) of globular clusters within the area covered by our observations can be roughly estimated as about 550 clusters, by extrapolating their luminosity function over the whole range of $T_1$ magnitudes. Thus, the number of clusters that should be distributed within a circular area of radius 80\\arcmin around \\object{NGC 1399} may be calculated, just taking into account the ratio of the areas, as several $10^4$ clusters. For comparison, the total number of globular clusters associated with galaxies in the Fornax Cluster may also be roughly estimated as follows. The blue magnitude of all the galaxies included in the \\citet{fer89} catalogue is $B = 8.3$ mag; adopting for them a mean color index $B-V \\approx 0.8$ mag and the distance modulus mentioned above, we obtain an absolute visual magnitude $M_V = -23.8$. If we assume a \"typical\" specific frequency $S_\\mathrm{N}=5$ we conclude that about $1.6\\times 10^4$ globular clusters should be associated with galaxies in Fornax.  It is interesting to note that the number of globular clusters that we found within a circular area of radius 80\\arcmin around \\object{NGC 1399} is of the same order o larger than the estimated number of globulars associated to galaxies in the whole Fornax Cluster.   It is also likely that some globular clusters might have escaped from its parent galaxies and, after that, remained within the potential well of the Fornax Cluster as a whole \\citep[see, for instance,][]{kis99}.  \\citet{whi87} propose that the distribution of the stripped globular clusters within the cluster will follow the same density profile as the galaxies and they might form a kind of envelope around the central galaxy. Alternative origins for the intraclusters are mentioned by \\citet{wes95}, who speculate that they might have formed ``in situ\", without a parent galaxy, or during mergers of sub-systems with a high gas content.  Deeper images are required to clarify this picture and a new survey in the Fornax Cluster is in progress. The true nature of these candidates might be confirmed by means of spectra."
    },
    "0212/gr-qc0212105.txt": {
        "abstract": "s{ Recent developments on the rotational instabilities of relativistic stars are reviewed. The article provides an account of the theory of stellar instabilities with emphasis on the rotational ones. Special attention is being paid to the study of these instabilities in the general relativistic regime. Issues such as the existence relativistic r-modes, the existence of a continuous spectrum and the CFS instability of the w-modes are discussed in the second half of the article.}  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "introduction": "The oscillations and instabilities of relativistic stars gained a lot of interest in the last decades because of the possible detection of their associated gravitational waves. It is not impossible that every compact star in a specific period of its life will undergo an oscillatory phase, in which it becomes unstable. Only the dynamical instabilities were thought to be a relevant source of gravitational waves, whereas instabilities due to dissipation mechanisms were believed to be only of academic interest. This belief has been dramatically altered during the last five years after it was discovered that for a specific class of rotational perturbations, the so called r-modes\\cite{Nils98,FM98}, the instability due to gravitational radiation\\cite{Chandra70a,Chandra70b,FS78a,FS78b} has the potential of being a prime source for gravitational waves. It was subsequently shown that this instability has many interesting astrophysical implications, which attracted the attention of both relativists and astrophysicists.  In the first section of this review we shortly describe the various stellar instabilities and how they operate, for a more detailed review one can refer to a recent article by Stergioulas\\cite{Nick98}. The second section is devoted to the r-mode instability with the focus on its application to gravitational-wave research and astrophysics, for more details see a recent review by Andersson and Kokkotas\\cite{AK01}. The last section is devoted to an analysis of new features brought about by the relativistic treatment of the problem. Questions on the existence relativistic r-modes and a continuous spectrum are discussed in detail. A new result that is presented is the existence of the w-mode instability in ultra-compact stars, i.e.~the instability due to the existence of an ergosphere of the spacetime or w-modes\\cite{KRA02}.   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": ""
    },
    "0212/astro-ph0212368_arXiv.txt": {
        "abstract": "We present optical and infrared spectra of SN~1999ex, which are characterized by the lack of strong hydrogen lines, weak optical He I lines, and strong He I $\\lambda$10830,20581. SN~1999ex provides a clear example of an intermediate case between pure Ib and Ic supernovae, which suggests a continuous spectroscopic sequence between SNe~Ic to SNe~Ib. Our $UBVRIz$ photometric observations of SN~1999ex started only one day after explosion, which permitted us to witness an elusive transient cooling phase that lasted 4 days. The initial cooling and subsequent heating due to $^{56}$Ni$\\rightarrow$$^{56}$Co$\\rightarrow$$^{56}$Fe produced a dip in the lightcurve which is consistent with explosion models involving core collapse of evolved massive helium stars, and not consistent with lightcurves resulting from the thermonuclear runaway of compact white dwarfs. ",
        "introduction": "In a rare occurrence the spiral galaxy IC~5179 ($cz$=3,498 $km$ $s^{-1}$) produced two supernovae (SNe) in an interval of only three weeks. The first object (SN~1999ee) was a Type Ia event discovered by us 10 days before maximum \\cite{maza99}. The early discovery motivated us to use the YALO and 0.9-m telescopes at the Cerro Tololo Inter-American Observatory in order to secure nightly $UBVRIz$ photometric observations of this event, and the YALO and Las Campanas 1-m and 2.5-m telescopes to obtain $JHK$ photometry. Although the second object (SN~1999ex) exploded three weeks later and was promptly present in our CCD images we did not notice its presence. Its discovery had to await independent observations obtained at Perth Observatory \\cite{martin99}. Once SN~1999ex was reported to the IAU Circulars we initiated an optical and infrared (IR) spectroscopic followup using the European Southern Observatory  NTT and Danish 1.5-m telescope at La Silla, and the VLT at Cerro Paranal. Our spectroscopic and photometric observations of SN~1999ex constitute an unprecedented dataset which provides support to our understanding of the nature of core collapse SNe. In this paper we show some of our observations and their interpretation. For a detailed report of our observations the reader is referred to \\cite{hamuy02}, \\cite{stritzinger02}, and \\cite{krisciunas02}. ",
        "conclusions": ""
    },
    "0212/astro-ph0212474_arXiv.txt": {
        "abstract": "{Optical spectral variability of quasars and BL Lac Objects is compared by means of the spectral variability parameter $\\beta$ \\citep{tre02}. Both kinds of objects change their spectral slopes $\\alpha$, becoming bluer when brighter, but BL Lac Objects have smaller $\\beta$ values and are clearly separated from quasars in the $\\alpha-\\beta$ plane. Models accounting for the origin of the variability are discussed for both classes of objects. ",
        "introduction": "Variability of the spectral energy distribution (SED) of Active Galactic Nuclei (AGNs) is a powerful tool to investigate the role of the main emission processes in different AGN classes, and the origin of their variations. The most common behavior in the optical band is that AGNs become bluer, i.e. their spectrum becomes harder, when brighter. This has been shown for individual Quasi Stellar Objects (QSOs) and Seyferts \\citep{cut85,ede90,kin91,pal94} and for one complete sample, i.e. for the 42 PG QSOs monitored by \\citet{giv99} in B and R for 7 years. Evidence based on two epochs has been found also for the faint QSO sample in the SA 57 \\citep{tre01}. The same trend is apparently shared by BL Lac Objects, as shown by \\citet{dam02}, who present 5-year long B, V, R, I light curves for eight objects. A quantitative estimate of color variability is necessary to compare the observational data with emission models and to compare different AGN classes. In a previous paper \\citep{tre02}, we introduced the spectral variability parameter $\\beta \\equiv \\Delta \\alpha / \\Delta \\log f_{\\nu}$, $f_{\\nu}$ being the specific flux and $\\alpha\\equiv\\partial\\log f_{\\nu}/\\partial\\log\\nu$ the spectral slope. From the average $\\alpha$ and $\\beta$ values, it was possible to derive constraints on the variability mechanisms. The spectral variability parameter was then estimated \\citep{vag02} for the BL Lac Objects of \\citet{dam02}. Observations and models can be compared in the $\\alpha-\\beta$ plane, which is reported in Fig. 1 for both the PG QSOs and the BL Lac Objects. ",
        "conclusions": "We show that the spectral variability parameter $\\beta$ is a powerful tool to discriminate between different models of the variability of AGNs. Hot spots on the disk, likely produced by local instabilities, are able to account for the observed spectral variability of QSOs.  We show that  BL Lacs clearly differ from QSOs in their $\\alpha,\\beta$ distribution. A simple model representing the variability of a synchrotron component can account for the observed $\\alpha$ and $\\beta$ values.  In the framework of wide field variability studies, we stress that observations in at least two photometric bands, repeated on the same field at many epochs, would allow a detailed test of variability models, extending our knowledge of the emission processes in AGNs."
    },
    "0212/astro-ph0212197_arXiv.txt": {
        "abstract": "We present results from a {\\it Chandra} observation of the NGC\\,346 cluster, which is the ionizing source of N66, the most luminous \\hii\\ region and the largest star formation region in the SMC. In the first part of this investigation, we have analysed the X-ray properties of the cluster itself and the remarkable star \\hd. But the field contains additional objects of interest. In total, 75 X-ray point sources were detected in the {\\it Chandra} observation: this is five times the number of sources detected by previous X-ray surveys. We investigate here their characteristics in detail. Due to high foreground absorption, the sources possess rather high hardness ratios. Their cumulative luminosity function appears generally steeper than that for the rest of the SMC at higher luminosities. Their absorption columns suggest that most of the sources belong to NGC\\,346. Using DSS data and new $UBVRI$ imaging with the ESO 2.2m telescope, we also discovered possible counterparts for 32 of these X-ray sources and estimated a B spectral type for a large number of these counterparts. This tends to suggest that most of the X-ray sources in the field are in fact X-ray binaries. Finally, some objects show X-ray and/or optical variability, with a need for further monitoring. ",
        "introduction": "The launch of the {\\it Chandra} satellite provides an opportunity to explore the X-ray sky with a far greater sensitivity and spatial resolution than ever before. In a given field of interest, these characteristics enable the discovery of numerous X-ray sources in addition to the main target(s). Although often regarded as secondary, these sources provide important information which can improve our knowledge of the X-ray emission mechanisms in the Universe.\\\\  These serendipitous discoveries also enable us to study the source distribution in galaxies, especially Magellanic dwarf galaxies in our case. Most of the dwarf galaxies studied are at a distance of a few Mpc (e.g. Martin, Kobulnicky, \\& Heckman 2002), and their X-ray observations only sample the population with L$_X$ $>$ 10$^{36}$ erg s$^{-1}$. On the other hand, the relative closeness of the Magellanic Clouds enable us to probe the X-ray sources with luminosities as low as 10$^{32}$ erg s$^{-1}$. Putting together the informations on faint and bright X-ray sources will ultimately lead to a better understanding of this type of galaxies.\\\\  We have obtained a deep {\\it Chandra} observation of the giant \\hii\\ region N66 \\citep{he56}, the largest star formation region in the Small Magellanic Cloud (SMC). The large number of massive stars and the presence of the remarkable star \\hd\\ make the NGC\\,346 field one of the best opportunities to conduct an investigation of the X-ray domain. \\\\  In the first part of this analysis (Naz\\'e et al. 2002, paper I), we have presented the characteristics of the cluster, the star \\hd\\ and its close neighborhood. The cluster itself is relatively faint and most of its emission seems correlated with the location of the brightest stars in the core. However, the level of X-ray emission could not be explained solely by the emission from individual stars. The {\\it Chandra} observation also provides the first X-ray detection of \\hd. In X-rays, the star, that underwent a LBV-type eruption in 1994, appears very bright, comparable only to the brightest WR stars in the Galaxy. This high luminosity could be explained either by colliding winds in the binary system or by post-eruption effects. Finally, a bright, extended X-ray emission seems to surround this star. It is probably due to a SNR which may or may not be related to \\hd\\ itself (see paper I). \\\\  In this second paper, we will focus on the other X-ray sources present in the field. First, we will describe in \\S~2 the observations used in this study. The detected sources, their hardness ratios (HRs), and their spectral characteristics will then be discussed in \\S~3, 4, and 5, respectively. Next, we will describe the overall properties of the point sources' population in \\S~6, present their possible counterparts in \\S~7 and investigate their variability in \\S~8. Finally, we will give a summary in \\S~9.\\\\ ",
        "conclusions": "In this series of articles, we report the analysis of the {\\it Chandra} data of N66, the largest star formation region of the SMC. In this second article, we have focused on the other sources detected in the field. The X-ray properties of 75 point sources, of which 32 may possess an optical counterpart, have been analysed. Their cumulative luminosity function is steeper than the global one of the SMC at higher luminosities. Using new photometry of the NGC\\,346 field, we estimate that a large number of the counterparts are B-type stars, suggesting that many of the X-ray sources may be X-ray binaries. Considering their absorption column and the photometry of their counterparts, we also conclude that most of these X-ray sources probably belong to NGC\\,346. Finally, due to their variability, some of the objects should be monitored in the future, both in the visible and X-ray wavebands.\\\\"
    },
    "0212/astro-ph0212312_arXiv.txt": {
        "abstract": "{ We apply the 2-dimensional high-resolution density field of galaxies of the Early Data Release of the Sloan Digital Sky Survey with a smoothing lengths 0.8~\\Mpc\\ to extract clusters and groups of galaxies, and a low-resolution field with smoothing lengths 10~\\Mpc\\ to extract superclusters of galaxies.  We investigate properties of density field clusters and superclusters and compare properties of these clusters and superclusters with Abell clusters, and superclusters found on the basis of Abell clusters.  We found that clusters in high-density environment have a luminosity a factor of $\\sim 5$ higher than in low-density environment.  There exists a large anisotropy between the SDSS Northern and Southern sample in the properties of clusters and superclusters: most luminous clusters and superclusters in the Northern sample are a factor of 2 more luminous than the respective systems in the Southern sample. ",
        "introduction": "Clusters and groups of galaxies are the basic building blocks of the Universe on cosmological scales.  The first catalogues of clusters of galaxies (Abell \\cite{abell}, Zwicky et al. \\cite{zwicky}) were constructed by visual inspection of the Palomar Observatory Sky Survey plates.  More recent catalogues of clusters, as well as catalogues of groups of galaxies, have been derived using catalogues of galaxies (\\cite{1983ApJS...52...89H}, \\cite{1997MNRAS.289..263D}). Moving up the hierarchy of large-scale structure, galaxy cluster catalogues themselves have been used to define still larger systems such as superclusters of galaxies (Einasto et al. \\cite{e1994},~\\cite{e1997}, ~\\cite{e2001}, hereafter E94, E97 and E01).  The goal of the present paper is to map the Universe up to redshift $z=0.2$ and to find galaxy clusters and superclusters using the density field method.  The application of the density field is not new.  In the pioneering study by \\cite{1982ApJ...254..437D} the density field was used to calculate the gravitational field of the nearby Universe.  Gott, Melott \\& Dickinson ~(\\cite{gmd86}) used the density field to investigate topological properties of the Universe. \\cite{1990ApJ...364..370B} calculated the potential, velocity, and density fields from redshift-distance data.  Saunders et al. (\\cite{s91}) applied the density field to map the Universe, to find superclusters and voids, and calculated moments of the density field. Marinoni et al. (\\cite{mar99}) reconstructed real space local density field. In these studies nearby optical or infrared galaxy samples were used.  More recently, Hoyle et al. (\\cite{h2002}) used smoothed 2-dimensional density fields from volume limited subsamples of SDSS EDR galaxies to discuss the 2-dimensional geometry of the large-scale matter distribution in comparison with $\\Lambda$CDM simulations, and Sheth et al. (\\cite{s2002}) advertised a new method of evaluating isodensity contours of smoothed 3-dimensional density fields from simulations for characterizing topological properties of the supercluster-void network.  \\begin{table*} \\caption[]{Data on SDSS EDR galaxies, clusters and superclusters}  \\label{Tab1} \\[ \\begin{tabular}{cccccccccccc} \\hline \\noalign{\\smallskip} Sample& DEC & RA & $\\Delta$RA & $\\alpha_1$&$M_1^{\\ast}$&$\\alpha_2$&$M_2^{\\ast}$& $N_{\\rm gal}$&$N_{\\rm cl}$&$N_{\\rm ACO}$& $N_{\\rm scl}$\\\\  \\noalign{\\smallskip} \\hline \\noalign{\\smallskip}  SDSS.N& $0^{\\circ}$&$190.25^{\\circ}$&$90.5^{\\circ}$ & $-1.06$&$-21.55$&$-1.22$&$-20.80$& 15209& 2868&22&24 \\\\ SDSS.S& $0^{\\circ}$&$23.25^{\\circ}$&$65.5^{\\circ}$& $-1.06$&$-21.40$&$-1.10$&$-20.71$&11882&2287&16&16\\\\ \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\] \\end{table*}    We calculate the density field of the Sloan Digital Sky Survey Early Data Release (SDSS EDR) by Stoughton et al. (\\cite{s02}), as described by H\\\"utsi et al. (\\cite{hytsi02}, hereafter Paper I), to find clusters and superclusters of galaxies, and to investigate their properties.  Clusters of galaxies from the SDSS were extracted previously by \\cite{2002PASJ...54..515G} using the cut and enhance method; \\cite{2002AJ....123...20K} compared various cluster detection algorithms based on SDSS data.  In this paper we define clusters as enhancements of the density field and use various smoothing lengths to separate systems of galaxies of different size and luminosity.  We use a high-resolution density field to find clusters and groups of galaxies.  For simplicity, we use the term ``DF-clusters'' for both groups and clusters found in the high-resolution density field of galaxies.  Similarly, we use a low-resolution density field to construct a catalogue of superclusters of galaxies, and denote them as ``DF-superclusters''. DF-clusters and superclusters are defined as enhancements of the density field, DF-clusters in a fixed volume ($\\pm 2.5$~\\Mpc\\ from the centre), and DF-superclusters as high-density regions surrounded by a fixed isodensity contour.  In determining DF-clusters and superclusters we take into account known selection effects.  We shall investigate some properties of DF-clusters and superclusters, and study these clusters and superclusters as tracers of the structure of the local Universe.   This study is of exploratory character to find the potential of SDSS data to analyse the structure of the Universe both on small and large scales.  In this stage we use the fact that SDSS EDR covers only relatively thin slices; thus we calculate the density field in 2 dimensions only.  As more data will be made available we plan to use a full 3-dimensional data set to detect clusters and superclusters of galaxies.  In this exploratory stage of the study we will make no attempt to convert distances of galaxies and systems of galaxies from redshift space to true space. This correction applies only to positions of galaxies and clusters, not to their luminosity.  Due to smoothing of the density field, small-scale corrections (e.g., the ``Finger-of-God'' Effect) are practically flattened out.  Although there are large-scale corrections due to the apparent contraction of superclusters in redshift space (the Kaiser Effect, \\cite{1984ApJ...284L...9K}), the actual positional shifts of supercluster centres are very small; furthermore, they do not alter cluster and galaxy positions in tangential direction.  The only large-scale redshift distortion is the radial contraction of superclusters, not the number and luminosity of clusters within superclusters.  Since supercluster shapes are not the main target of this present study, we may ignore this effect for now.  In Section 2 we give an overview of the observational data. In Section 3 we find DF-clusters and investigate their properties.  Similarly, in Section 4 we compose a catalogue of DF-superclusters, identify them with conventional superclusters, and study their properties.  In Section 5 we continue the study of DF-clusters and superclusters, derive the luminosity function of DF-clusters, and analyse these systems as tracers of the large-scale structure of the universe. Section 6 brings our conclusions.  The three-dimensional distribution of DF-clusters and superclusters in comparison with Abell clusters and superclusters is shown on the web-site of Tartu Observatory. The analysis of the density field of LCRS shall be published by Einasto et al. (\\cite{e02b}). ",
        "conclusions": "Results of this paper and Paper I are of a methodical and quantitative character.  Methodical aspects concern the application of the density field method with various smoothing lengths to display and describe systems of galaxies of various scales.  We calculated density field for a number of smoothing lengths from 0.8 to 16~\\Mpc. In this way systems of galaxies on various scales could be visualised and their mutual relationship could be studied.  The high-resolution density field, with dispersion 0.8~\\Mpc\\, was used to define density field clusters and to investigate the structure of superclusters.  The low-resolution density field, with dispersion 10~\\Mpc\\, was used to define superclusters of galaxies and to study global properties of clusters and superclusters.  The inspection of the distribution of clusters in superclusters and voids suggests that galaxy systems have in both regions similar shape in that the dominant structural elements are single or multi-branching filaments.  Massive superclusters have dominantly a multi-branching morphology; less massive superclusters have various morphologies, including compact, filamentary, multi-branching, and diffuse systems.  Quantitative results concern the luminosity function of galaxies and properties of clusters and superclusters.  We have found in Paper I and confirmed in this paper that it is impossible to find a global set of parameters of the luminosity function which can be applied in all cases.  We found that parameters of the luminosity function are different for high- and low-density regions: galaxies in high-density regions are more luminous.  This result confirms earlier findings by Lindner et al. (\\cite{l95}): bright galaxies define larger voids than faint ones. More recently this conclusion was obtained by \\cite{1998ApJ...505...25B}, \\cite{2000ApJ...545....6B}, and by \\cite{2001MNRAS.328...64N}.  The present study indicates that the effect is larger than previously suspected: luminosities of the brightest galaxies in high-density regions exceed  the luminosities of the brightest galaxies in low-density regions by a factor of about 5.  The values of the parameters of the luminosity function are also different for nearby and distant parts of the survey and for the Northern and Southern slices.  A distance dependence of parameters of the luminosity function has been found for deeper surveys spanning redshift interval up to $z=1$ by \\cite{1999ApJ...518..533L} and by \\cite{2001ApJ...560...72S} in the Canadian Network for Observational Cosmology Redshift Survey.  However, in this survey the distance dependence appears only by comparison of nearby ($z \\sim0$) and distant ($z > 0.2$) parts of the survey.  Thus it seems improbable that this effect can explain our results on the distance dependence in a relatively small redshift interval $0 < z \\leq 0.2$.  Parameters of the luminosity function depend also on the density of the environment.  This dependence influences properties of the density field. High-density regions contain brighter galaxies than do low-density ones.  This difference leads to different selection effects for galaxies in high- and low-density environments.  In a high-density environment, due to selection effects, faint galaxies in clusters cannot be observed, but clusters themselves are visible (since they contain at least one galaxy bright enough) and the total luminosity of clusters can be estimated to take into account unobserved galaxies.  In a low-density environment all galaxies of the cluster may be too faint, and the cluster may not be seen at all. In another words, distance-dependent selection effects influence clusters and superclusters in different ways.    The density field was calculated using the expected total luminosities of clusters of galaxies, including the expected luminosities of galaxies too faint or too bright to be included in the redshift survey.  The correction for unobserved galaxies was made assuming a Schechter luminosity function for galaxies. Our analysis shows that it is impossible to correct the density field so that general properties of the density field and properties of clusters of galaxies are correct for the same set of parameters of the galaxy luminosity function. For this reason we have used two sets of parameters of the luminosity function. Parameter set 1 has a strong bright-end (parameter $M^{\\star} = -21.55$) and corresponds to galaxies in a high-density environment which dominate clusters observed at large distance; this set yields correct properties of clusters of galaxies, but does not include faint clusters at large distances, and thus gives too low of a density for distant superclusters. Parameter set 2 was obtained for the whole region under study, and it has moderate bright-end (parameter $M^{\\star} = -20.80$). In this parameter set ALL faint invisible galaxies at large distance are included within visible clusters, including galaxies that belong to non-detected clusters; thus clusters themselves become too luminous with increasing distance, but supercluster properties are correct.   Comparing clusters in different environments we have found that there exists a strong dependence of cluster properties on the density of the large-scale environment: clusters located in high-density environments are a factor of $5 \\pm 2$ more luminous than clusters in low-density environments.  Finally we found that there exists a large difference between properties of clusters and superclusters in the Northern and Southern slices of the SDSS EDR survey: clusters and superclusters in the Northern slice are more luminous than those in the Southern slice by a factor of 2.  This difference may be due to differences in the location of slices with respect to the very large-scale environment. If this conclusion is confirmed by future observations one must conclude that the formation and evolution of galaxies and systems of galaxies of various scales depends on the nearby as well as on the large-scale environment.  Richer superclusters have more luminous galaxies and clusters.  On smaller scales this tendency has been observed as a difference between properties of clusters within superclusters and in voids.  Now we see that a similar difference may be observed on much larger scales."
    },
    "0212/astro-ph0212124_arXiv.txt": {
        "abstract": "Bose-Einstein condensation (BEC) and evaporation of transverse photons from the Bose condensate is studied in the case when the density of plasma does not change. The generation of the longitudinal photons (photonikos) by the transverse photons (photons) in terms of the Cherenkov type of radiation in a uniform plasma is demonstrated. The Bogoliubov energy spectrum is derived for photonikos.  A new physical phenomena of the \"Compton\" scattering type in nonlinear photon gas is discussed. To this end, a new version of the Pauli equation in the wavevector representation is derived. The formation of Bose-Einstein condensate and evaporation of photons from the condensate is investigated by means of the newly derived Fokker-Planck equation for photonikos. The relevance of this work to recent discovery of black hole X-ray jets is pointed out. ",
        "introduction": "An immense amount of research has been carried out on propagation of relativistically intense electromagnetic (EM) waves into an isotropic plasma demonstrating that the relativistic oscillatory motion of electrons causes a whole set of interesting and salient phenomena \\cite{shu}-\\cite{kiv}, relevant to the study of laser accelerators of electrons, ions and photons, laser fusion, nonlinear optics, etc. Some of them have already been confirmed by experiments due to the recent progress in compact, high-power, short pulse laser technology. In our previous paper \\cite{tsin} we have suggested a new mechanism for ultrahigh gradient electron and ion acceleration. In addition we have shown over $10^7G$ generated magnetic fields by the nonpotential ponderomotive force.  The above treatments were restricted to the case of monochromatic EM waves. However, the interaction of relativistically intense radiation with a plasma leads to several kinds of instabilities and the initially coherent spectrum may eventually broaden, and naturally for ultrashort pulses the initial bandwidth is increasingly broad. Moreover, in astronomical plasmas there are a variety of sources of radiation, and in this case we speak about the average density of radiation from all the sources and their spectral distribution. Hence, the natural state of the strong radiation of EM field is with a broad spectrum. In order to study the interaction of spectrally  broad and relativistically intense EM waves with a plasma, it was necessary to derive a general equation for the EM spectral intensity. Such an investigation was reported recently \\cite{ltsin98}, \\cite{ntsin98}, where the authors derived a general kinetic equation for the photons in a plasma.  It should be emphasized that the radiation can be in two distinct states. Namely, one is when the total number of photons is not conserved. A good example is a black-body radiation. Another situation is when the total number of photons is conserved. Both these states have been studied in several aspects, mostly for the weak radiation.  It is well known that the nature of EM waves in a vacuum is quit different from the one in a medium. In the vacuum EM wave exists only in motion, however the light can be stopped (wavevector $\\vec{k}=0,\\ \\omega=\\omega(0)\\neq 0$) in different mediums and wave-guides. As was shown in Ref.\\cite{ltsin96} the photons acquire the rest mass and become one of the Bosons in plasmas and posses all characteristics of nonzero rest mass, i.e., we may say that the photon is the elementary particle of the optical field.  Besides, a several reviews and books have been published on theory of Bose-Einstein condensation (BEC) in a quantum Bose liquid \\cite{lif} and in trapped gases \\cite{dal},\\cite{ant}. Kompaneets \\cite{kom} has shown that the establishment of equilibrium between the photons and the electrons is possible through the Compton effect. In his consideration, since the free electron does not absorb and emit, but only scatters the photon, the total number of photons is conserved. Using the kinetic equation of Kompaneets, Zel'dovich and Levich \\cite{zel} have shown that in the absence of absorption the photons undergo BEC. Such a possibility of the BEC occurs in the case, when the processes of change of energy and momentum in scattering dominates over the processes involving change of the photon number in their emission and absorption.  Very recently it was shown that exists an another new mechanism of the creation of equilibrium state and Bose-Einstein condensation in a nonideal dense photon gas \\cite{ltsin02}. More importantly a new effect was predicted in the same paper, namely that the inhomogeneous dense photon gas can be found in the intermediate state.  In the present paper, we consider the BEC and evaporation of the transverse photons (photons) from the Bose condensate. In our study we assume that the intensity of radiation (strong and super-strong laser pulse, non-thermal equilibrium cosmic field radiation, etc.) is sufficiently large, so that the photon-photon interaction can become more likely than the photon-electron interaction. We will show that for certain conditions the variation of the plasma density can be neglected in comparison with the variation of the photon density. In such case the elementary excitations represent the longitudinal photons (photonikos), for which we will derive the well known Bogoliubov's energy spectrum.  The paper is organized as follows. First in Sec.II a basic equations, describing the relativistic photon-plasma interactions, is presented. The problem of stability of the photon flow is discussed in Sec.III. The derivation of the Bogoliubov energy spectrum for the photonikos is given in the same section. Then in Sec.IV, we derive the Pauli kinetic equation from the Wigner-Moyal equation. Section V is devoted to BEC. The Fokker-Planck equation for photonikos, by which we discuss the possibility of the creation of Bose-Einstein condensate and evaporation of the photons from the condensate, is obtained in the same section. Finally, a brief summary and discussion of our results are given in the last section. ",
        "conclusions": "We have investigated a class of problems involving the interaction of spectrally  broad and relativistically intense EM radiation with a plasma in the case when the photon-photon interaction dominates the photon-particle interactions. We have presented a new concept of the establishment of equilibrium between the photon and the wavepacket of EM field (the dense photon bunch). This is a fully relativistic effect due to the strong EM radiation. We have established the condition under which the variation of the plasma density can be neglected in comparison with the variation of the photon density. In such case the elementary excitations represent the photonikos, for which we have derived the well known Bogoliubov energy spectrum. We have studied the BEC and evaporation of the photons from the Bose condensate in the case when the density of plasma does not change. To this end, from the Wigner-Moyal-Tsintsadze equation \\cite{ltsin98},\\cite{ntsin98},\\cite{men} we have derived a new version of the Pauli kinetic equation for the photon gas. For the case, when the wavevector and frequency of the photoniko is small in comparison with the  wavevector and frequency of the photon, we have derived the Fokker-Planck equation for photonikos. We have presented a simple model, which exhibits the possibility of the creation of Bose-Einstein condensate and evaporation of photons from the condensate. We think that these processes can be detectable in next generation experiments with appropriate instrumentation. In fact, a number of experiments have been carried out in which plasmas are irradiated by laser beams with intensities up to $3\\cdot 10^{20}W/cm^2$. At such intensities the photon density is of the order of $n_\\gamma\\sim 10^{29}cm^{-3}.$ For the plasma densities up to $n_e\\sim 10^{19}$, we should expect that the Bose-Einstein condensate becomes observable. The theory developed in this paper should also be the case for astrophysical objects, such as a black hole. In this connection we speculate that the recently observed radiation from the black hole may be attributed to the evaporated photons from the Bose condensate. That is initially all photons fall into Bose-Einstein condensate, and then after a certain time, as discussed in this paper, some photons undergo evaporation from the condensate. These processes may also explain the observable variation in radiation intensity, from being undetectable to one of the brightest sources on the sky. Note that these sources can turn off for decades, and new ones are always being found. In addition Universe is filled with a gas of photonikos. This gas may play the decisive role in the expansion of the Universe. Moreover, it may also be useful for explaining certain processes in supernovae explosion. Finally, the present theory may also find a valuable application in the future space technology, as well as in nonlinear optics."
    },
    "0212/astro-ph0212254_arXiv.txt": {
        "abstract": "s{I discuss recent developments in the field of relativistic jets in AGNs. After a brief review of our current knowledge of emission from Blazars, I discuss some consequences of the recent detection made by {\\it Chandra} of X-ray emission from extended jets. Finally I report some recent results on the problem of the connection between accretion and jets, study that in principle could shed light on the important issue of jet formation.} ",
        "introduction": "Jets are ``pipelines'' through which energy and matter originating from the central region of AGNs can flow out from the galaxy and, as in FRII radio galaxies, reach the radio-lobes (e.g. Begelman, Blandford \\& Rees 1984). The comprehension of such structures represents a great challenge for Astrophysics: the ultimate goals include the understanding of the physical processes able to produce such energetics and collimated outlows and the complex dynamics of the jet and its interaction with the environment.  In the last decade, thanks to new instruments (in particular the $\\gamma $-ray telescope EGRET, the TeV Cherenkov telescopes and {\\it Chandra}), this field has received a great impulse. In the following I briefly present our current view of the principal physical processes acting in jets, from the innermost portion of the jet close to the central ``engine'' to the external regions. Finally I discuss some works in progress regarding the connection between the inner portion of the jet and the outer regions. ",
        "conclusions": ""
    },
    "0212/astro-ph0212062_arXiv.txt": {
        "abstract": "We present an equation of state and radiative opacities for a strongly magnetized hydrogen plasma at magnetic fields $B$, temperatures $T$, and densities $\\rho$ typical for atmospheres of isolated neutron stars. The first- and second-order thermodynamic functions, monochromatic radiative opacities, and Rosseland mean opacities are calculated and tabulated, taking account of partial ionization, for $8\\times10^{11}$ G $\\leq B\\leq 3\\times10^{13}$ G, $2\\times10^5$ K $\\leq T\\leq 10^7$ K, and a wide range of $\\rho$. We show that bound-bound and bound-free transitions give an important contribution to the opacities at $T \\la (1$---$5)\\times 10^6$ K in the considered range of $B$ in the outer neutron-star atmosphere layers, which may substantially modify the X-ray spectrum of a typical magnetized neutron star. In addition, we re-evaluate opacities due to free-free transitions, taking into account the motion of both interacting particles, electron and proton, in a strong magnetic field. Compared to the previous neutron-star atmosphere models, the free-free absorption is strongly suppressed at photon frequencies below the proton cyclotron frequency. The latter result holds for any field strength, which prompts a revision of existing theoretical models of X-ray spectra of magnetar atmospheres. ",
        "introduction": "\\label{sect-intro} Models of neutron star atmospheres are needed for interpretation of their spectra and cooling. These atmospheres differ from the atmospheres of ordinary stars because of the high gravity and magnetic fields (for review, see, e.g., \\citealp{Pavlov95,elounda}).  A magnetic field is called \\emph{strong}\\ if the electron cyclotron energy $\\hbar\\omc=\\hbar eB/\\mel c$ exceeds 1 a.u. -- i.e., the field strength $B$ is higher than $B_0=\\mel^2 c\\, e^3/\\hbar^3 = 2.3505\\times10^9$~G, where $\\mel$ is the electron mass, $e$ the elementary charge, and $c$ the speed of light. Usually the field is called \\emph{superstrong} if $\\hbar\\omc > \\mel c^2$, that is $B > B_\\mathrm{r}=\\mel^2 c^3/e\\hbar = 4.414\\times10^{13}$~G. Most of the radio pulsars have magnetic fields $B \\sim 10^{12}$---$10^{13}$ G \\citep*{tml93}, whereas anomalous X-ray pulsars and soft gamma repeaters are thought to have superstrong fields (e.g., \\citealp{mereghetti01,thompson00}, and references therein). Non-negligible amount of neutral atoms can exist in the photosphere at typical neutron-star temperatures $T \\sim 10^6$ K (\\citealp*{PCS99}, hereafter Paper I). A strong magnetic field enhances atomic binding and makes the quantum-mechanical characteristics of an atom dependent on its motion across the field (see \\citealp{Lai01} for a recent review). In photospheres of the neutron stars, the field is, as a rule, \\emph{strongly quantizing}, i.e., it sets all the electrons on the ground Landau level. This occurs if $\\beta_\\mathrm{e} \\gg 1$ and $\\rho<\\rho_B$, where \\beq \\beta_\\mathrm{e} = \\hbar\\omc/\\kB T \\approx 134.3\\,B_{12}/T_6 , \\label{beta-e} \\eeq $\\rho$ is the density, and $\\rho_B = \\mH/(\\pi^2\\sqrt{2}\\,\\am^3) \\approx 7100\\,B_{12}^{3/2}$ \\gcc\\ (for the hydrogen plasma). Here and hereafter, $\\mH = \\mpr+\\mel$, $\\mpr$ is the proton mass,  $\\am=(\\hbar c/eB)^{1/2}$ is the \\emph{magnetic length}, $\\kB$ is the Boltzmann constant, $B_{12}=B/10^{12}$~G, and $T_6=T/10^6$~K.  Opacities for the two polarization modes of radiation are quite different in strongly magnetized plasmas (e.g., \\citealp{Pavlov95}, and references therein), which makes thermal emission of neutron stars polarized and anisotropic \\citep{Zavlin95}. The mean opacities are strongly reduced at $\\beta_\\mathrm{e}\\gg1$ (e.g., \\citealp{silyak}); thus the bottom of the photosphere is shifted to high densities (e.g., \\citealp{Pavlov95,LS97}).  The chemical composition of neutron-star atmospheres is not precisely known. Just after the neutron star birth in a supernova explosion, the outer stellar envelope is most probably composed of iron. However, light elements may be brought to the surface later (e.g., by fallback, accretion, or encounters with comets). Because of rapid gravitational sedimentation, the lightest element will cover the surface (see \\citealp*{Brown02}). About $10^{12}$---$10^{14}$ grams of hydrogen ($<10^{-19}$ M$_\\odot$) is sufficient to fill the entire photosphere.  \\citet{Shib92} presented the first model of hydrogen atmospheres with strong magnetic fields. Later it was developed beyond the diffusion approximation \\citep{SZ95} and used for astrophysical predictions (e.g., \\citealp{Zavlin95}; \\citealp*{Zane00,Zane01,HoLai,HoLai02,LaiHo02,Ozel01,Ozel02}) and for interpretation of observed neutron-star spectra (e.g., \\citealp*{PSZ95,PSZ96,Pavlov95,Ozel-ea}).  The above studies assume that the atmosphere is fully ionized. Meanwhile, it was recognized long ago (e.g., \\citealp{Miller92}) that a significant contribution to the opacities of neutron-star photospheres with strong magnetic fields might come from bound-bound and bound-free absorption by atoms. Examples of monochromatic opacities in partially ionized iron \\citep*{RRM} and hydrogen \\citep*{PCS00} atmospheres confirmed this conjecture. In Paper I we have presented an equation of state (EOS) of a partially ionized hydrogen plasma for the values of $T$ and $B$ typical for atmospheres of the radio pulsars. Here we report results of extensive calculations of thermodynamic functions based on the theory developed in Paper~I, supplemented by calculations of the opacities (monochromatic and Rosseland mean). Partial ionization and plasma nonideality are taken into account for $11.9 \\leq \\log_{10} B/\\mathrm{G} \\leq 13.5$ and $5.3 \\leq \\log_{10} T/\\mathrm{K} \\leq 7.0$. Bound-bound and bound-free radiative transitions are treated within the framework of a previously developed theory \\citep{PP95,PP97}. The free-free absorption cross sections are re-evaluated. Whereas the previous authors considered photoabsorption by an electron scattered off a fixed Coulomb center, we take into account the finite proton mass, which has a nontrivial effect on the photoabsorption in a quantizing magnetic field.  The paper is composed as follows. In Sect.\\ \\ref{sect-input} we formulate the main assumptions and give the basic formulae used in our work. Section \\ref{sect-EOS} presents the EOS of partially ionized hydrogen under conditions in neutron-star photospheres. In Sect.\\ \\ref{sect-cross} we discuss various contributions to the hydrogen photoabsorption cross sections in strong magnetic fields and derive a new formula for the free-free cross section. Opacities of hydrogen photospheres of the neutron stars are discussed in Sect.\\ \\ref{sect-opac}. Appendices give some detail of calculation of the free-free cross sections. ",
        "conclusions": "We have calculated the EOS and radiative opacities of fully and partially ionized hydrogen plasmas in a wide range of densities, temperatures, and magnetic fields typical for photospheres of the strongly magnetized neutron stars. The first- and second-order thermodynamic functions, non-ionized fractions, and effective Rosseland mean opacities are published in the electronic form.  The opacities are calculated more accurately than in the previous publications. In particular, we take into account suppression of the free-free absorption below the proton cyclotron frequency, which was overlooked previously. This effect reduces the opacities of the ionized component of the plasma by orders of magnitude at photon energies $\\hbar\\omega\\lesssim0.3\\,\\hbar\\omp\\sim 0.02\\,B_{12}$ keV, which necessitates a revision of the previously published models of X-ray spectra of magnetars \\citep{Zane01,HoLai,HoLai02,Ozel01,Ozel02}. On the other hand, the bound-bound and bound-free absorption, neglected in the previous models of neutron-star atmospheres, increase the opacities by more than one order of magnitude at $\\hbar\\omega\\sim(0.1$--3) keV in the outer atmosphere layers of the ordinary neutron stars with $B\\sim10^{12}$--$10^{13.5}$ G and $T < (1$--$3)\\times10^6$ K, which can also significantly affect the spectra.  One can expect that the effect of the bound species on the EOS and opacities is as important for magnetars (despite their supposedly higher temperatures) as for the ordinary neutron stars. To check this, we need to extend our model to higher $B$; preliminary high-$B$ results \\citep{CDP02} support this anticipation."
    },
    "0212/astro-ph0212581_arXiv.txt": {
        "abstract": "We construct models of triaxial galactic nuclei containing central black holes using the method of orbital superposition, then verify their stability by advancing $N$-body realizations of the models forward in time. We assume a power-law form for the stellar density, $\\rho \\propto r^{-\\gamma}$, with $\\gamma=1$ and $\\gamma=2$; these values correspond approximately to the nuclear density profiles of bright and faint galaxies respectively. Equidensity surfaces are ellipsoids with fixed axis ratios. The central black hole is represented by a Newtonian point mass. We consider three triaxial shapes for each value of $\\gamma$: almost prolate, almost oblate and maximally triaxial. Two kinds of orbital solution are attempted for each mass model: the first including only regular orbits, the second including chaotic orbits as well. We find that stable configurations exist, for both values of $\\gamma$, in the maximally triaxial and nearly-oblate cases; however steady-state solutions in the nearly-prolate geometry could not be found. A large fraction of the mass, of order 50\\% or more, could be assigned to the chaotic orbits without inducing evolution. Our results demonstrate that triaxiality may persist even within the sphere of influence of the central black hole, and that chaotic orbits may constitute an important building block of galactic nuclei. ",
        "introduction": "Schwarzschild (1979) demonstrated how to construct self-consistent models of stellar systems in the absence of analytic expressions for the orbital integrals. His method consists of three steps. i) Represent the stellar system by a smooth density law and divide it into discrete cells; ii) compute a library of orbits in the potential corresponding to the assumed density law, and record the time spent by each orbit in the cells; iii) find a linear combination of orbits that reproduces the cell masses. Using his method, Schwarzschild (1979, 1982) demonstrated self-consistency of triaxial mass models with and without figure rotation. Most of the orbits in his solutions were regular, i.e. non-chaotic (Merritt 1980). Subsequently, Statler (1987) found a variety of self-consistent solutions for the integrable, or ``perfect,'' triaxial mass models in which all orbits are regular.  Models like these with large, constant-density cores are now known to be poor representations of elliptical galaxies, almost all of which have high central densities (\\cite{cra93}; Ferrarese et al. 1994). Stellar densities rise toward the center approximately as power laws, $\\rho \\propto r^{-\\gamma}$. Fainter galaxies have steeper cusps, $\\gamma \\approx 2$, while brighter galaxies have weaker cusps, $0\\lap\\gamma\\lap 1$, and exhibit an obvious break in the surface brightness profile. Following this discovery, Schwarzschild (1993) investigated triaxial models with singular density profiles, $\\rho\\sim r^{-2}$, and Merritt \\& Fridman (1996) constructed self-consistent solutions for triaxial galaxies with both weak and strong central cusps. A significant portion of the phase space in these models was found to be occupied by stochastic orbits. Furthermore, triaxial self-consistency could sometimes only be achieved by including some stochastic orbits. Models containing stochastic orbits can represent bona-fide equilibria as long as the stochastic orbits are represented as fully-mixed ensembles (Merritt \\& Fridman 1996; Merritt \\& Valluri 1996).  In the last decade, evidence has grown that supermassive black holes are generic components of galactic nuclei (Ho 1999). There are roughly a dozen galaxies in which a compelling case for the presence of a supermassive black hole can be made based on the kinematics of stars or gas (Merritt \\& Ferrarese 2001), as well as a number of active galactic nuclei in which the kinematics of the broad emission line region imply the existence of a supermassive black hole (Peterson 2002). Inferred masses range from $\\sim 10^{6}\\msun$ to $\\sim 10^{9.5}\\msun$ and correlate well with stellar velocity dispersions (e.g. \\cite{fer01}) and bulge luminosities (e.g. \\cite{mcd02}).  The possibility of maintaining triaxiality within a galactic nucleus containing a supermassive black hole remains a topic of interest. Very close to the black hole, the gravitational force can be considered a perturbation to the Kepler problem and the phase space is essentially regular (\\cite{mev99}; \\cite{sas00}; \\cite{pom01}, hereafter Paper I). Farther from the black hole, the fraction of chaotic orbits increases, up to a radius where the enclosed stellar mass is a few times the black hole mass; beyond this radius essentially all centrophilic orbits are chaotic (Paper I). The tube orbits remain mostly regular since they avoid the destabilizing center. The persistence of regular orbits throughout the region where the gravitational force from the black hole dominates, leaves open the possibility of constructing self-consistent solutions. Furthermore there is growing observational evidence for the existence of bar-like distortions at the very centers of galaxies (e. g. \\cite{ers02}).  In Paper II (Poon \\& Merritt 2002), we presented preliminary results showing that self-consistent and stable triaxial equilibria could be constructed for power-law nuclei with certain axis ratios. In this paper, we present a more detailed investigation of triaxial black-hole nuclei. We find that stationary solutions are possible only for certain shapes; mass models that are too near to prolate axisymmetry always evolve toward axisymmetry.  The properties of the mass models are presented in \\S2. Orbital solutions for various shapes and density profiles are presented in \\S3 and their stability is tested by N-body simulation in \\S4 and \\S5. Limits on the chaotic mass fraction are discussed in \\S6. \\S7 discusses some implications for the nuclear dynamics of galaxies. ",
        "conclusions": "We have shown that long-lived triaxial configurations are possible for nuclei containing black holes. Models with $T=0.5$ (maximally triaxial) and $T=0.25$ (oblate/triaxial) were constructed and found to be stable, retaining their non-axisymmetric shapes until the end of the integration interval, equal to several crossing times. Models with $T=0.75$ (prolate/triaxial) were always found to evolve rapidly to axisymmetry; we speculate that prolate/triaxial nuclei do not exist. The evolution seen in the nearly-prolate models does not appear to be a consequence of orbital chaos; indeed, in our stable solutions, we were able to replace a surprisingly large fraction of the regular orbits by chaotic orbits without inducing noticeable evolution in their shapes. We found that at least 50\\%, and perhaps as much as 75\\%, of the mass could be placed on chaotic orbits in the maximally triaxial and oblate/triaxial solutions. Such models violate Jeans's theorem in its standard form (e.g. \\cite{bt87}) but are consistent with a generalized Jeans's theorem (\\cite{mer99}) if we assume that the chaotic building blocks are ``fully mixed,'' that is, that they approximate a uniform population of the accessible phase space. This appears to be the case for the chaotic orbits in our models which have a very short mixing time. While a sudden onset of chaos can effectively destroy triaxiality in models containing a large population of regular box orbits (\\cite{meq98}; \\cite{sel01}), our work shows that at least the central parts of galaxies containing black holes can remain triaxial even when dominated by chaotic orbits.  Our results have possibly important implications for the rate at which stars are fed to supermassive black holes in galactic nuclei. In spherical or axisymmetric nuclei, the feeding rate is determined by the rate at which stars on eccentric orbits are scattered into the loss cone, the phase-space region defined by orbits with pericenters lying within the black hole's tidal disruption radius. In the case of chaotic orbits in a triaxial nucleus, each passage brings the star near to the center, and the time required for a star to pass within a distance $r_t$ of the black hole should scale roughly as $r_t^{-1}$ (e.g. \\cite{geb85}). Thus even in the absence of gravitational scattering, the loss cone would remain full and the feeding rate could be orders of magnitude higher than in axisymmetric nuclei. We will examine these ``chaotic loss cones'' in detail in an upcoming paper (Merritt \\& Poon 2003).  While our results strengthen the case for triaxiality in galactic nuclei, the case for nuclear triaxiality could be made even more compelling by the detection of isophotal twists or minor-axis rotation at the very centers of galaxies. Such observations will be challenging, requiring two-dimensional data on an angular scale that resolves the black hole's sphere of influence. Existing integral-field spectrographs on ground-based telescopes (e.g. SAURON, \\cite{bac01}) can only achieve this resolution for the nearest galaxies. Equally valuable would be $N$-body studies demonstrating that triaxial nuclei can form in realistic mergers.  M.Y. Poon would like to thank Andrew Mack for stimulating discussions and constructive comments. This work was supported by NSF grants AST 96-17088 and AST 00-71099 and by NASA grants NAG5-6037 and NAG5-9046. M.Y. Poon is grateful to the Croucher Foundation for a postdoctoral fellowship.  \\clearpage"
    },
    "0212/hep-ph0212018_arXiv.txt": {
        "abstract": "We have earlier shown that cosmic strings moving through the plasma at the time of a first order quark-hadron transition in the early universe can generate large scale baryon inhomogeneities. In this paper, we calculate detailed structure of these inhomogeneities at the quark-hadron transition. Our calculations show that the inhomogeneities generated by cosmic string wakes  can strongly affect nucleosynthesis calculations. A comparison with observational data suggests that such baryon inhomogeneities should not have existed at the nucleosynthesis epoch. If this disagreement holds with more accurate observations, then it will lead to the conclusions that cosmic string formation scales above $10^{14} - 10^{15}$ GeV may not be consistent with nucleosynthesis and CMBR observations. Alternatively, some other input in our calculation should be constrained, for example, if the average string velocity remains sufficiently small so that significant density perturbations are never produced at the QCD scale, or if strings move ultra-relativistically so that string wakes are very thin, trapping negligible amount of baryons. Finally, if quark-hadron transition is not of first order then our calculations  do not apply. ",
        "introduction": "Recent measurements of the cosmic microwave background radiation (CMBR) anisotropy have reached a sufficiently high level of precision that stringent bounds can be put on various cosmological parameters such as baryon to entropy ratio $\\eta$. It is certainly quite remarkable that the calculations of standard big-bang nucleosynthesis (SBBN) are reasonably consistent with these stringent bounds on $\\eta$. Though several modifications to SBBN are still being considered to better account for the abundances of light elements. One such possibility discussed extensively in the literature is the so called inhomogeneous big bang nucleosynthesis (IBBN) \\cite{ibbn1,ibbn2} where nucleosynthesis takes place in the presence of baryon number inhomogeneities. Various models have been worked out where inhomogeneities of a particular shape and size are taken and their effects on nucleosynthesis are calculated. Different values of the light elemental abundances are calculated and compared with the observed values \\cite{ibbn2}. These calculations are supplemented with the investigations of baryon inhomogeneity generation during a first order quark-hadron phase transition  \\cite{bfluct,bfluct2}. Though, it is  fair to say that current observations do not support any strong deviation from the SBBN calculations. Calculations of IBBN, such as those in ref. \\cite{ibbn1,ibbn2}, therefore, can be used to constrain the presence of baryon fluctuations in the early universe.  In a previous paper \\cite{layek}, we had shown that baryon inhomogeneities on large-scales will be generated by cosmic string wakes during the quark-hadron transition. This arises due to the fact that when there are density fluctuations present in the universe, (for example, those which eventually lead to the formation of structure we  see today), then resulting temperature fluctuations, even if they are of small magnitude, can affect the dynamics of a first order phase transition in crucial ways. There have been many discussions of the effects of inhomogeneities on the dynamics of a first order quark-hadron transition in the universe \\cite{impur,inhm}. For example, Christiansen and Madsen have discussed \\cite{impur} heterogeneous nucleation of hadronic bubbles due to presence of impurities. Hadronic bubbles  are expected to nucleate at these impurities with enhanced rates. Recently, Ignatius and Schwarz have proposed \\cite{inhm} that the presence of density fluctuations (those arising from inflation) at  quark-hadron transition will lead to splitting of the region in hot and cold regions with cold regions converting to hadronic phase first. Baryons will then get trapped in the (initially) hotter regions. Estimates of sizes and separations of such density fluctuations were made in ref.\\cite{inhm} using COBE measurements of the temperature fluctuations in CMBR. In ref. \\cite{layek}, we considered the effect of cosmic string induced density fluctuations on quark-hadron transition and showed that it can lead to formation of extended planar regions of baryon inhomogeneity.  We mention here again, as discussed in ref.\\cite{layek}, that there has been extensive study of density fluctuations generated by cosmic strings from the point of view of structure formation \\cite{str1}, and it is reasonably clear that recent measurements of temperature anisotropies in the microwave background by BOOMERANG, and MAXIMA experiments \\cite{expt} at angular scales of $\\ell \\simeq$ 200 disfavor models of structure formation based exclusively on cosmic strings \\cite{str2,str3}. However, even with present models of cosmic string network evolution, it is not ruled out that cosmic strings may contribute to some part in the structure formation in the universe. Further, cosmic strings generically arise in many Grand Unified Theory (GUT) models. If the GUT scale is somewhat lower than $10^{16}$ GeV then the resulting cosmic strings will not be relevant for structure formation (for a discussion of these issues, see \\cite{str3}). However, they may still affect various stages of the evolution of the universe in important ways. Our study in ref.\\cite{layek} (see, also ref.\\cite{ew}), and the present study are motivated from this point of view.  In this paper we determine the detailed structure of the baryon inhomogeneities created by the cosmic string wakes \\cite{layek}. We find that the magnitude and length scale of these inhomogeneities are such that they should survive until the stage of nucleosynthesis, affecting the calculations of abundances of light elements.  A comparison with observational data suggests that such baryon inhomogeneities should not have existed at the nucleosynthesis epoch. If this disagreement holds with more accurate observations then it will lead to the conclusion  that cosmic string formation scales above $10^{14} - 10^{15}$ GeV may not be consistent with nucleosynthesis and CMBR observations. Alternatively, some other input in our calculation should be constrained, for example, if the average string velocity remains sufficiently small so that significant density perturbations are never produced at the QCD scale, or if strings move ultra-relativistically so that string wakes are very thin, trapping negligible amount of baryons.  Of course entire discussion of this paper is applicable only when quark-hadron transition is of first order.  The paper is organized in the following manner. In section II, we briefly discuss the nature of density fluctuations as expected from cosmic strings moving through a relativistic fluid. In section III we discuss the dynamics of quark-hadron transition in the presence of string wakes, and discuss how baryons are concentrated in sheet like regions inside these wakes. Section IV discusses the results of our calculations where we present the detailed structure of the baryon inhomogeneities. In section V we discuss the effects, these baryon inhomogeneities surviving until the nucleosynthesis stage, can have on the abundances of light elements, and discuss the constraints on the cosmic string models arising from observations of various abundances. Conclusions are presented in section VI. ",
        "conclusions": "We have calculated the detailed structure of the baryon inhomogeneities created by the cosmic string wakes \\cite{layek}. We find that the magnitude and length scale of these inhomogeneities is such that they survive until the stage of nucleosynthesis, affecting the calculations of abundances of light elements. A comparison with observational data suggests that such baryon inhomogeneities should not have existed at the nucleosynthesis epoch. If this disagreement holds with more detailed calculations and more accurate observations, then it will lead to the conclusion that cosmic string formation scales above a value of about $10^{14} -10^{15}$ GeV are not consistent with nucleosynthesis and CMBR observations. Alternatively, some other input in our calculation should be constrained, for example, the average string velocity can be sufficiently small so that significant density perturbations are never produced at the QCD scale, or strings may move ultra-relativistically so that resulting wakes are very thin, and trap a negligible amount of baryon number. Finally, all these considerations are valid only when quark-hadron transition is of first order.  There are many uncertainties in our model, for example treatment of multiple wakes is rather ad hoc. Similarly, we have tried to use results from ref.\\cite{ibbn2} adopting them for our case even though detailed geometry of baryon inhomogeneity in our case is different. A more careful,  detailed calculation of abundances of elements is needed for the present case.  The uncertainties in various observations of abundances of elements, as well as CMBR anisotropy will be reduced as precision of various measurements gets better. Then one will be able to say with a greater certainty whether IBBN results puts a strong restriction on the density fluctuations, and hence on cosmic string parameters, or the order of quark-hadron phase transition.  \\vskip .2in \\centerline {\\bf ACKNOWLEDGMENTS} \\vskip .1in  We are very thankful to Rajiv Gavai and Rajarshi Ray for many useful suggestions and comments. We also thank Mark Trodden for very useful discussions and suggestions."
    },
    "0212/astro-ph0212188_arXiv.txt": {
        "abstract": "We present results of hydrodynamical simulations of young supernova remnants. To model the ejecta, we  use several models (discussed in literature) of type Ia supernova explosions with different abundances. Our hydro models are one-dimensional and spherically symmetrical, but they take into account ionization kinetics with all important processes. We include detailed calculations for the X-ray emission, allowing for time-dependent ionization and recombination. In particular, we compare the computed X-ray spectra with recent XMM-Newton observations of the Tycho SN remnant. Our goal is to find the most viable thermonuclear SN model that gives good fits to both these X-ray observations and typical SN~Ia light curves. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/astro-ph0212141_arXiv.txt": {
        "abstract": "We present a method for the study of second order superhorizon perturbations in multi field inflationary models with non trivial kinetic terms. We utilise a change of coordinates in field space to separate isocurvature and adiabatic perturbations generalizing previous results. We also construct second order gauge invariant variables related to them. It is found that with an arbitrary metric in field space the isocurvature perturbation sources the gravitational potential on long wavelengths even for ``straight'' trajectories. The potential decouples from the isocurvature perturbations if the background fields' trajectory is a geodesic in field space. Taking nonlinear effects into account shows that, in general, the two types of perturbations couple to each other. This is an outline of a possible procedure to study nonlinear and non-Gaussian effects during multifield inflation. ",
        "introduction": "Cosmological perturbation theory is a rather arcane subject. The reason is that in a general perturbed spacetime there is no privileged coordinate system with respect to which one can define perturbations. So perturbations can change when we change the coordinates. The study of general relativistic perturbations was pioneered in \\cite{lifs} and studied by many authors since (see e.g \\cite{mfb} for a comprehensive review). Let us briefly recall a more formal presentation of what is usually meant when one talks of perturbations in general relativity \\cite{j.m.s}. One considers a five dimensional space composed of the background spacetime ${\\mathcal{M}}_0$ and, stacked above it, perturbed spacetimes ${\\mathcal{M}}_{\\epsilon}$ parametrized by the parameter $\\epsilon$. We implicitly assume some sort of differentiable structure on this 5-D space such that these perturbed spacetimes can be considered ``close'' to ${\\mathcal{M}}_0$. On these spacetimes live tensor fields $T$. One then defines a vector field $X$, the integral curves of which are used for identifying points on ${\\mathcal{M}}_{\\epsilon}$ with points on the background ${\\mathcal{M}}_0$. The choice of $X$ is completely arbitrary and is called a choice of {\\it{gauge}}. In general, any tensor $T$ can be expanded as a taylor series \\beq T_{0}+\\delta_{X}T=\\phi^{*}_{X\\epsilon}(T_{\\epsilon})=T_0+\\epsilon\\pounds_XT\\mid_0+\\frac{1}{2}\\epsilon^2\\pounds_X\\pounds_XT\\mid_0+...\\,, \\eeq or, calling the various terms the perturbations at various orders, \\beq \\phi^{*}_{X\\epsilon}(T_{\\epsilon})=T_0+{\\delta}T^{(1)}+{\\delta}T^{(2)}+...\\, , \\eeq where $\\phi^{*}_{X\\epsilon}$ is the pullback along $X$ on the background manifold ${\\mathcal{M}}_0$ of a tensor that lives on a perturbed spacetime $\\mathcal{M}_{\\epsilon}$ parameter distance $\\epsilon$ away from the background. Hence the vector field $X$ alows us to define perturbations in a meaningful way. The choice of another vector field  $Y=(1,Y^{\\mu})$ defines a different gauge and one finds that perturbations differ when defined in different gauges: \\beq \\delta_YT-\\delta_XT=\\epsilon\\pounds_{\\xi^{(1)}}T\\mid_0+\\frac{1}{2}\\epsilon^2\\pounds_{\\xi^{(1)}}^2T\\mid_0+\\epsilon^2\\pounds_{\\xi^{(1)}}\\delta^{(1)}_XT+\\frac{1}{2}\\epsilon^2\\pounds_{\\xi^{(2)}}T\\mid_0+...\\,, \\eeq where $\\delta^{(1)}_XT\\equiv\\epsilon\\pounds_XT\\mid_0$ is the linear perturbation of $T$ in the ``$X$ gauge'' and $\\xi^{(1)}\\equiv{Y-X}$, $\\xi^{(2)}\\equiv{[X,Y]}$ are vector fields which lie on ${\\mathcal{M}}_0$ and are independent of each other. Hence, by a suitable choice of ${\\xi^{(1)}}$ and ${\\xi^{(2)}}$, a gauge condition can be imposed order by order.  Expansion (3) also suggests a strategy for identifying gauge invariant quantities at second order. Observe that the first and fourth terms on th r.h.s of (3) are essentially the same (the transformations they define have the same functional form). This means that any linear combination $f(\\delta{T}^{(1)})$ of first order variables which is gauge invariant to first order will also be gauge invariant w.r.t that part of the second order transformation which corresponds to the $\\frac{1}{2}\\epsilon^2\\pounds_{\\xi^{(2)}}T\\mid_0$ term in (3). The remaining terms, $\\frac{1}{2}\\epsilon^2\\pounds_{\\xi^{(1)}}^2T\\mid_0+\\epsilon^2\\pounds_{\\xi^{(1)}}\\delta^{(1)}_XT$, are all composed of products of first order quantities. So in seeking gauge invariant combinations at second order we must look for appropriate quadratic terms of first order quantities that will cancel these quadratic terms in (3). If this can be done in a unique way then the form of a gauge invariant quantity at first order will dictate its form at second order.  We will now give an explicit example of the construction of a second order gauge invariant variable corresponding to the well known first order gauge invariant quantity (first introduced in \\cite{luk}, see also \\cite{mfb}) \\beq \\cR=\\psi+\\frac{\\ch}{\\varphi_{0}'}\\delta\\varphi . \\eeq In general every quantity will be expanded in orders like in (2). For example \\bea g_{\\mu\\nu}&=&g^{(0)}_{\\mu\\nu}+{\\delta}g^{(1)}_{\\mu\\nu}+\\frac{1}{2}{\\delta}g^{(2)}_{\\mu\\nu}+...\\nonumber\\\\ \\varphi&=&\\varphi^{(0)}+{\\delta}\\varphi^{(1)}+\\frac{1}{2}{\\delta}\\varphi^{(2)}+... \\eea e.t.c. In particular, writing the general perturbed metric element as \\beq ds^2=a^2[(1+2\\phi)d\\tau^2-2B_{i}dx^id\\tau-[(1-2\\psi)\\gamma_{ij}+2E_{ij}]dx^idx^j] \\eeq we have \\bea g_{00}&=&a(\\tau)^2\\left(1+2\\phi^{(1)}+\\phi^{(2)}+...\\,\\right)\\\\ g_{0i}&=&a(\\tau)^2\\left(B_i^{(1)}+\\frac{1}{2}B_i^{(2)}+...\\,\\right)\\\\ g_{ij}&=&-a(\\tau)^2\\left[\\left(1-2\\psi^{(1)}-\\psi^{(2)}+...\\right)\\delta_{ij}+2E_{ij}^{(1)}+E_{ij}^{(2)}+...\\,\\right]. \\eea Then, from eqn. (3) one can calculate the formulae for the gauge transformations of the relevant quantities. For an extensive account of second order gauge transformations, explicit formulae and some specific examples the reader can see \\cite{bruni,bruni2} and references therein.  In general, formulae for perturbations at second order can be complicated and calculations rather tedious. In this paper we will make a number of simplifying assumptions. We will ignore vectors (hence spatially indexed quantities are given by derivatives of scalars) and, mainly, we will drop terms containing more than one spatial gradients. They are expected to be unimportant on scales longer than the hubble radious. Although there is no rigorous justification for the latter approximation it is expected to capture the main affects on superhorizon scales \\cite{long}. Within such an approach, initial conditions at horizon crossing can be set by linear theory. Then, the long wavelength equations can be used to calculate the nonlinearities induced during the superhorizon evolution. The authors of \\cite{uzan} adopted such a procedure and showed that it is possible for significant nongaussianities to be generated in the adiabatic mode from the long wavelength evolution in multifield inflationary models. Their calculation ignored metric perturbations which are included here (see next section).  With these approximations in mind we have the following formulae for a gauge transformation: At first order \\bea \\tilde\\phi_{(1)}&=&\\phi_{(1)}+{\\xi^0_{(1)}}'+\\ch\\xi_{(1)}^0,\\\\ \\tilde{B}_{(1)i}&=&{B}_{(1)i}-\\partial_i\\xi^0_{(1)}+\\xi_{i(1)}',\\\\ \\tilde\\psi_{(1)}&=&\\psi_{(1)}-\\ch\\xi_{(1)}^0,\\\\ \\tilde{E}_{(1)ij}&=&{E(1)}_{ij},\\\\ \\tilde\\delta\\varphi_{(1)}&=&\\delta\\varphi_{(1)}+\\varphi'\\xi^0_{(1)}, \\eea and at second order \\bea \\tilde\\phi_{(2)}=\\phi_{(2)}+{\\xi^0_{(2)}}'+\\ch\\xi_{(2)}^0 &+&\\xi_{(1)}^0\\left[2\\left(\\phi_{(1)}'+2\\ch\\phi_{(1)}\\right)+{\\xi_{(1)}^0}''+5\\ch{\\xi_{(1)}^0}'+\\left(\\ch'+2\\ch^2\\right)\\xi_{(1)}^0\\right]\\nonumber\\\\ &+&2{\\xi_{(1)}^0}'\\left(2\\phi_{(1)}+{\\xi_{(1)}^0}'\\right), \\eea \\bea \\tilde{B}_{(2)i}&=&{B}_{(2)i}-4\\phi_{(1)}\\partial_i\\xi_{(1)}^0+\\xi_{(1)}^0\\left[2\\left({B_i^{(1)}}'+2\\ch B_i^{(1)}\\right)-\\partial_i{\\xi^0_{(1)}}'+{\\xi_i^{(1)}}''-4\\ch\\left(\\partial_i\\xi_{(1)}^0-{\\xi_i^{(1)}}'\\right)\\right]\\nonumber\\\\ &+&{\\xi^0_{(1)}}'\\left(2B_i^{(1)}-3\\partial_i\\xi^0_{(1)}+{\\xi_{i(1)}}'\\right)+{\\xi_{(1)}^j}'\\left(-4\\psi_{(1)}\\delta_{ij}+2E_{(1)ij}\\right)-\\partial_i\\xi^0_{(2)}+\\xi_{i(2)}', \\eea \\beq \\tilde\\psi_{(2)}=\\psi_{(2)}-\\ch\\xi^0_{(2)}+\\xi_{(1)}^0\\left[2\\left(\\psi_{(1)}'+2\\ch\\psi_{(1)}\\right)-\\left(\\ch'+2\\ch^2\\right)\\xi_{(1)}^0-\\ch{\\xi_{(1)}^0}'\\right], \\eeq \\beq \\tilde{E}_{(2)ij}={E}_{(2)ij}+2\\xi^0_{(1)}\\left({E}_{(1)ij}'+2\\ch{E}_{(1)ij}\\right) \\eeq \\beq \\tilde\\delta\\varphi_{(2)}=\\delta\\varphi_{(2)}+\\varphi'\\xi_{(2)}^0+\\xi_{(1)}^0\\left(\\varphi''\\xi_{(1)}^0+\\varphi'{\\xi_{(1)}^0}'+2\\delta\\varphi_{(1)}'\\right). \\eeq Note that, as mentioned before, the part of the transformations in (15) - (19) containing the vector field $\\xi_{(2)}$ is exactly the same as the first order case, eqn's (10) - (14). From the above we see that the variable (4) at second order transforms like \\bea \\frac{\\ch}{\\varphi_0'}\\delta\\tilde\\varphi^{(2)}+\\tilde\\psi^{(2)}=\\frac{\\ch}{\\varphi_0'}\\delta\\varphi^{(2)}+\\psi^{(2)}&+&\\left[\\ch\\frac{\\varphi_0''}{\\varphi_0'}-\\left(\\ch'+2\\ch^2\\right)\\right](\\xi_{(1)}^0)^2\\nonumber\\\\ &+&2\\left(\\frac{\\ch}{\\varphi_0'}\\delta\\varphi_{(1)}'+\\psi_{(1)}'+2\\ch\\psi_{(1)}\\right)\\xi_{(1)}^0 \\eea As expected the transformation contains only products of first order quantities. Therefore we seek to construct a gauge invariant quantity at second order by adding a quadratic combination of first order quantities that will transform appropriately. By inspection we see that it must contain $\\psi$, $\\delta\\varphi'$ and $\\psi'$ and it must not contain $\\delta\\varphi$. So we must have \\beq \\left(A\\psi+C\\psi'+D\\delta\\varphi'\\right)\\left(E\\psi+G\\psi'+H\\delta\\varphi'\\right). \\eeq Noting that \\bea \\psi'\\,\\rightarrow\\,\\psi'-\\ch'\\xi^0-\\ch{\\xi^{0}}'\\\\ \\delta\\varphi'\\,\\rightarrow\\,\\delta\\varphi'+\\varphi_0''\\xi^0+\\varphi_0'{\\xi^{0}}' \\eea and that we must not have terms involving ${\\xi^{0}}'$ we see that we have 2 options. We either set \\bea D=C\\frac{\\ch}{\\varphi_0'}\\\\ H=G\\frac{\\ch}{\\varphi_0'} \\eea which eliminates the terms involving ${\\xi^{0}}'$ in the transformation or take \\bea D=H=\\frac{4\\pi}{m_p^2}\\ch\\varphi_0'\\\\ C=G=\\ch^2-\\ch'. \\eea and use the background equations of motion. In both cases we are forced to consider $A=E$, $C=G$ and $D=H$ and we end up with the same variable \\cite{mat} \\bea \\cR_{(2)}&=&\\left[\\psi_{(2)}+\\frac{\\ch}{\\varphi_{0}'}\\delta\\varphi_{(2)}\\right]\\nonumber\\\\ &+&\\frac{\\left[\\psi_{(1)}'+2\\ch\\psi_{(1)}+\\frac{\\ch}{\\varphi_0'}\\delta_{(1)}\\varphi'\\right]^2}{\\ch'+2\\ch^2-\\ch\\frac{\\varphi_0''}{\\varphi_0'}} \\eea which is invariant under the transformations (17) and (19). ",
        "conclusions": "We have touched upon the issue of studying second order perturbations in multifield inflationary models and defined gauge invariant variables - equations (102) and (103) - on supehorizon scales also to second order. The latter can be constructed given the solution to lowest non linear order in a given gauge. We have presented a more geometrical method for the splitting of isocurvature and adiabatic perturbations and applied it to a two field  model with a diagonal but otherwise arbitrary metric. We showed that in this case naive ``straight'' trajectories do not lead to the decoupling of adiabatic from isocurvature perturbations. We identified the type of curves for which this happens. A perturbative approach by which one can study the nonlinear evolution of perturbations in such models to second order was suggested. The resulting equations have the same form as the first order linear ones but with new terms appearing on the right hand side. These new terms are quadratic in the first order perturbations and therefore they can be considered as known ``sources'' from the solution of the first order problem. Such a formalism can be used to calculate the amount of nongaussianity produced in such models by treating the perturbations as gaussian stochastic fields when they become superhorizon and then study their non linear evolution from that point. The resulting non gaussianity will be of the $\\chi^2$ type since we are considering quadratic products of gaussian fields. We have implicitly assumed a smoothing on scales larger than the horizon and dropped second order spatial gradients. Although straightforward to calculate, the resulting source terms, in principle known, are quite complicated. We will return to the issue of calculating non linear evolution in multifield inflationary models with a different approach in a future publication \\cite{gp}.   {\\bf Acknowledgements:} Many thanks to Carsten van de Bruck and Christopher Gordon for useful discussions and comments, Marco Bruni for bringing to my attention other works related to second order perturbation theory, and especially Paul Shellard for discussions and generous support.  \\appendix"
    },
    "0212/astro-ph0212007_arXiv.txt": {
        "abstract": "A number of `modified' Newtonian potentials of various forms are available in the literature which accurately approximate some general relativistic effects important for studying accretion discs around a Schwarzschild black hole. Such potentials may be called `pseudo-Schwarzschild' potentials because they nicely mimic the space-time around a non-rotating/slowly rotating compact object. In this paper, we examine the validity of the application of some of these potentials to study the spherically symmetric, transonic, hydrodynamic accretion onto a Schwarzschild black hole. By comparing the values of various dynamical and thermodynamic accretion parameters obtained for flows using these potentials with full general relativistic calculations, we have shown that though the potentials discussed in this paper were originally proposed to mimic the relativistic effects manifested in disc accretion, it is quite reasonable to use most of the potentials in studying various dynamical as well as thermodynamic quantities for spherical accretion to compromise between the ease of handling of a Newtonian description of gravity and the realistic situations described by complicated general relativistic calculations.  Also we have shown that depending on the chosen regions of parameter space spanned by specific energy ${\\cal E}$ and adiabatic index $\\gamma$ of the flow, one potential may have more importance than another and we could identify which potential is the best approximation for full general relativistic flow in Scwarzschild space-time for particular values of ${\\cal E}$ and $\\gamma$.   \\keywords {accretion, accretion discs --- black hole physics ---  hydrodynamics} ",
        "introduction": "Stationary, spherically symmetric and transonic hydrodynamic accretion of adiabatic fluid on to a gravitating astrophysical object at rest was studied in a seminal paper by Bondi (1952) using purely Newtonian Potential and by including the pressure effect of the accreting material. Later, Michel (1972) discussed fully general relativistic polytropic accretion on to a Schwarzschild black hole by formulating the governing equations for steady spherical flow of a perfect fluid in the Schwarzschild metric. Following Michel's relativistic generalization of Bondi's treatment, Begelman (1978) discussed some aspects of the critical (sonic) points of the flow for such an accretion. Using an unrelaxed mono-energetic particle distribution and assuming the fact that the relaxation time of such a particle distribution is very long compared to the typical flow time scale or dynamical time scale of steady accretion on to black holes, Blumenthal and Mathews (1976) developed a model where the connection between the nonrelativistic to the relativistic regime of the spherically accreting material could be established. Taking the fully ionized one-temperature ($T_{electron}=T_{proton}$) hydrogen gas (governed by an exact relativistic equation of state) to be the fundamental constituent of the accreting material, Brinkmann (1980) treated spherically symmetric stationary accretion in Schwarzschild space time and showed that the temperature of accreting material at the Schwarzschild radius is one order of magnitude smaller than the flow temperature obtained by using a simple polytropic equation of state. Recently, Malec (1999) provided the solution for general relativistic  spherical accretion with and without back reaction and showed that relativistic effects enhance mass accretion when back reaction  is neglected.\\\\ \\noindent Meanwhile, the theory of the accretion disc found prior importance because of the fact that disc accretion describes more realistic astrophysical situations found in nature. The beginning of modern accretion disc physics is traditionally attributed to the two classical articles by Shakura \\& Sunyaev (1973) and Novikov \\& Thorne(1973) . While Shakura \\& Sunyaev (1973) calculated the disk structure and related phenomena using purely Newtonian potential, Novikov \\& Thorne provided a fully general relativistic description of accretion discs around black holes; later on, some aspects of which were modified by Riffert and Herold(1995). However, rigorous investigation of transonic disc structure was found to be extremely complicated in full general relativistic space time (Chakrabarti (1996) and references therein). At the same time it was understood that as relativistic effects play an important role in the regions close to the accreting black hole (where most of the gravitational potential energy is released), purely Newtonian gravitational potential (in the form ${\\Phi}_{Newton}=-\\frac{GM}{r}$) cannot be a realistic choice to describe transonic black hole accretion in general. To compromise between the ease of handling of a Newtonian description of gravity and the realistic situations described by complicated general relativistic calculations, a series of `modified' Newtonian potentials have been introduced to describe the general relativistic effects that are most important for accretion disk structure around Schwarzschild and Kerr black holes ( see Artemova et. al.(1996) for further discussion). Introduction of such potentials allows one to investigate the complicated physical processes taking place in disc accretion in a semi-Newtonian framework by avoiding pure general relativistic calculations that most of the features of spacetime around a compact object are retained and some crucial properties of the analogous relativistic solutions of disc structure could be reproduced with high accuracy. Hence, those potentials might be designated as `pseudo-Kerr' or `pseudo- Schwarzschild' potentials, depending on whether they are used to mimic the space time around a rapidly rotating or non rotating/ slowly rotating (Kerr parameter $a\\sim0$) black hole respectively.\\\\ \\noindent It is important to note that although a number of such `pseudo' potentials are available in the literature to study various aspects of disc accretion, no such potentials are available which had been solely derived to describe spherically symmetric accretion on to a Schwarzschild (or Kerr) black hole. In this paper, we will concentrate on some of the `pseudo-Schwarzschild' {\\it disc} potentials (potentials introduced to study {\\it accretion discs} around a Schwarzschild black hole) to investigate whether those potentials could be used to study Spherical accretion, and if so, how `good' the choice would be for various such potentials. Also, we would like to check which potential among those would be the `best-fit' to approximate the full general relativistic description of transonic, spherically symmetric accretion on to a Schwarzschild black hole. In doing so, we solve the equations of motion of spherically accreting fluid in full Schwarzschild space-time as well as for motion under various `pseudo'-potentials, to study the variation of different dynamical and thermodynamic quantities (like Mach number of the flow, flow temperature  etc.) with radial distance measured from the accreting black hole for the {\\it full general relativistic spherical flow} (hereafter FGRSF) as well as for accretion using various `pseudo- Schwarzschild' potentials. We then compare the results obtained using such potentials with the solutions of exact relativistic problems in a Schwarzschild metric. The plan of the paper is as follows: In next section, we will describe four `pseudo-Schwarzschild' disc potentials available in the literature and some of their basic features. In \\S 3, we will provide the basic equations governing spherically symmetric accretion in full relativistic as well as in various `pseudo'--relativistic spacetimes.  In \\S 4 we will discuss  how to solve those equations to find various dynamical quantities which are to be mutually compared and we present our results. Finally in \\S 5 we conclude by discussing the suitability of various `pseudo' potentials in approximating the results obtained from exact relativistic calculations. For the rest of this paper, we will use the terms `modified- Newtonian potential' and `pseudo (Schwarzschild) potentials' synonymously. \\section {Some basic features of various `pseudo-Schwarzschild' potentials} From now, we will define the Schwarzschild radius $r_g$ as $$ r_g=\\frac{2G{M_{BH}}}{c^2} $$ (where  $M_{BH}$  is the mass of the black hole, $G$ is universal gravitational constant and $c$ is velocity of light in vacuum) so that the marginally bound circular orbit $r_b$ and the last stable circular orbit $r_s$ take the values $2r_g$ and $3r_g$ respectively for a typical Schwarzschild black hole. Also, total mechanical energy per unit mass on $r_s$ (sometimes called `efficiency' $e$) may be computed as $-0.057$ for this case. Also, we will use a simplified geometric unit throughout this paper where radial distance $r$ is be scaled in units of $r_g$, radial dynamical velocity $u$ and polytropic sound speed $a$ of the flow is scaled in units of $c$ (the velocity of light in vacuum), mass $m$ is scaled in units of $M_{BH}$ and all \\begin{figure} \\vbox{ \\vskip -5.8cm \\centerline{ \\psfig{file=MS1237f1.eps,height=15cm,width=15cm}}} \\noindent {{\\bf Fig. 1:} Newtonian potential $\\Phi_{Newton}(r)$ and other pseudo-potentials $\\Phi_i(r)$'s ($i=1,2,3,4$) are plotted as a function of the logarithmic radial distance from the accreting black hole. Except $\\Phi_2$, all other pseudo-potentials have a singularity at one Schwarzschild radius $r_g$. For $r~>~2r_g$, while the relative stiffness factor ${\\cal S}$ is maximum for $\\Phi_4$, it is minimum for $\\Phi_2$, see text for details.} \\end{figure} other derived quantities would be scaled accordingly. Also, for simplicity,we will use $G=c=M=1$. Below we would like to briefly describe four different `pseudo- Schwarzschild' potentials (expressed in the system of units discussed above) and to provide the `free fall' acceleration obtained using such potentials in compact form.\\\\ \\noindent Paczy\\'nski and Wiita (1980) introduced a `pseudo-schwarzschild' potential of the form $$ \\Phi_{1}=-\\frac{1}{2(r-1)} \\eqno{(1a)} $$ which accurately reproduces the positions of $r_s$ and $r_b$ and gives the value of efficiency to be $-0.0625$. Also the Keplarian distribution of angular momentum obtained using this potential is exactly same as that obtained in pure Schwarzschild geometry. Although this potential (as well as the other `pseudo' potentials available in the literature) does not satisfy the boundary condition exactly on the horizon, however, close to the horizon, accretion flows are supposed to be highly supersonic and the dynamical infall time scale becomes too small to allow significant radiation from infalling fluid. Even if there are significant amounts of radiation, it does not contribute too much to the overall radiation coming out from the disc, due to the fact that radiation emitted from the very close viscinity of the black hole is highly redshifted. Thus, errors in calculation close to the horizon may have very little impact on the emitted disc spectrum and for all practical purposes, the Paczy\\'nski and Wiita (1980) potential is considered to be the best approximation of pure Schwarzschild spacetime while treating the disc accretion (Artemova 1996) . It is worth mentioning here that this potential was first introduced to study a thick accretion disc with super Eddington Luminosity. Also, it is interesting to note that although it had been thought of in terms of disc accretion, it is spherically symmetric (with a scale shift of $r_g$). So in principle, it can definitely be used to study spherical accretion also on to a Schwarzschild black hole.\\\\ \\noindent To analyse the normal modes of accoustic oscillations within a thin accretion disc around a compact object (slowly rotating black hole or weakly magnetized neutron star), Nowak and Wagoner (1991) approximated some of the dominant relativistic effects of the accreting black hole (slowly rotating or nonrotating) via a modified Newtonian potential of the form $$ \\Phi_{2}=-\\frac{1}{2r}\\left[1-\\frac{3}{2r}+12{\\left(\\frac{1}{2r}\\right)}^2\\right ] \\eqno{(1b)} $$ $\\Phi_2$ has correct form of $r_s$ as in the Schwarzschild case but is unable to reproduce the value of $r_b$. This potential has the correct general relativistic value of the angular velocity (as measured at infinity) at $r_s$. Also it reproduces the radial epicyclic frequency $\\kappa$ (for $r>r_s$) close to its value obtained from general relativistic calculations. However, this potential gives the value of efficiency as $-0.064$ which is larger than that produced by $\\Phi_1$, hence the disc spectrum computed using $\\Phi_2$ would be more luminous compared to a disc structure studied using $\\Phi_1$.\\\\ \\noindent Remembering  that the free-fall acceleration plays a very crucial role in Newtonian gravity, Artemova et. al. (1996) proposed two different `pseudo' potentials to study disc accretion around a non-rotating black hole. The first potential proposed by them produces exactly the same value of the free-fall acceleration of a test particle at a given value of $r$ as is obtained for a test particle at rest with respect to the Schwarzschild reference frame, and is given by $$ \\Phi_{3}=-1+{\\left(1-\\frac{1}{r}\\right)}^{\\frac{1}{2}} \\eqno{(1c)} $$ The second one gives the value of the free fall acceleration that is equal to the value of the covariant component of the three dimensional free-fall acceleration vector of a test particle that is at rest in the Schwarzschild reference frame and is given by $$ \\Phi_{4}=\\frac{1}{2}ln{\\left(1-\\frac{1}{r}\\right)} \\eqno{(1d)} $$ Efficiencies produced by $\\Phi_3$ and $\\Phi_4$ are $-0.081$ and $-0.078$ respectively.The magnitude of efficiency produced by $\\Phi_3$ being maximum,calculation of disc structure using $\\Phi_3$ will give  the maximum amount of energy dissipation and the corresponding spectrum would be the most luminous one. However, as both $\\Phi_3$ and $\\Phi_4$ stems from the consideration of free fall acceleration and calculates the dependence of free fall accleration on radial distance in the Schwarzschild metric (which describes a spherically symmetric gravitational field in vacuum), it appears to be quite justified to use those potentials to study spherically symmetric accretion.\\\\ \\noindent From now we will refer to all these four potentials as $\\Phi_i$ in general where $\\left\\{i=1,2,3,4\\right\\}$ would correspond to $\\Phi_1$ (eqn. 1(a)), $\\Phi_2$ (eqn. 1(b)), $\\Phi_3$ (eqn. 1(c)) and $\\Phi_4$ (eqn. 1(d)) respectively. In Fig. 1, we plot various $\\Phi_i$'s as a function of the radial distance measured from the accreting black hole in units of $r_g$. Also in the same plot, purely Newtonian potential $\\Phi_{Newton}$ is plotted. If we now define a quantity ${\\cal S}_i$ to be the `relative stiffness' of a potential $\\Phi_i$ as: $$ {\\cal S}_i=\\frac{\\Phi_i}{r} $$ (that is, ${\\cal S}_i$ is a measure of the numerical value of any $i$th potential at a radial distance $r$), we find that for $r~>~2r_g$: $$ {\\cal S}_2~<~{\\cal S}_{Newton}~<~{\\cal S}_1~<~{\\cal S}_3~<~{\\cal S}_4 $$ which indicates that while $\\Phi_2$ is a `flatter' potential compared to the pure Newtonian potential $\\Phi_{Newton}$, all other `pseudo' potentials are `steeper' to  $\\Phi_{Newton}$ for $r~>~2r_g$. \\\\ \\noindent One can write the modulus of free fall acceleration obtained from all `pseudo' potentials except for $\\Phi_2$ in a compact form as $$ \\left|{{{{{\\Phi}^{'}}_{i}}}}\\right|=\\frac{1}{2{r^{2-{\\delta}_{i}} {\\left(r-1\\right)}^{\\delta_{i}}}} \\eqno{(2a)} $$ where ${\\delta_{1}}=2$, $\\delta_3=\\frac{1}{2}$ and $\\delta_4=1$. $\\left|{{{{{\\Phi}^{'}}_{i}}}}\\right|$ denotes the absolute value of the {\\it space derivative} of $\\Phi_i$, i.e., $$ \\left|{{{{{\\Phi}^{'}}_{i}}}}\\right|=\\left|{\\frac{d{\\Phi_i}}{dr}}\\right| $$ whereas acceleration produced by $\\Phi_2$ can be computed as, $$ {\\Phi_2}^{'}=\\frac{1}{2r^2}\\left(1-\\frac{3}{r}+\\frac{9}{2r^2}\\right) \\eqno{(2b)} $$ In the next section,we would like to describe how one can investigate transonic spherical accretion using these potentials. Also, we will discuss how to calculate various dynamical quantities for full general relativistic bondi flow in a Schwarzschild metric. One standard method to investigate classical transonic bondi flow is to formulate the basic conservation equations, i.e.,conservation of baryon number (obtained by integrating continuity equation) and conservation of specific energy (obtained by integrating Euler's equation using a specific equation of state for accreting matter) and then to simultaneously  solve these conservation equations to get critical (sonic) quantities as functions of various accretion parameters (like specific energy ${\\cal E}$, adiabatic index $\\gamma$ or  accretion rate ${\\dot M}_{in}$  in the flow) and also to calculate the values of various dynamical and thermodynamic quantities (like Mach number $M$, of the flow, flow temperature, $T$, etc.) as functions of various accretion parameters (like ${\\cal E}, \\gamma, {\\dot M}_{in}$ etc.) or radial distance (measured from the central accretor in the unit of Schwarzschild radius $r_g$). Also a common practice is to study the variation of Mach number $M$ (ratio of the local dynamical velocity $u$ to the local sound velocity $a$; $M=\\frac{u}{a}$) with radial distance $r$ (measured in units of $r_g$) to investigate the `transonicity' of the flow for a fixed set of input parameters. We will perform the above mentioned calculations for accretion in pure Schwarzschild spacetime as well as for accretion in `pseudo-schwarzschild' space time using $\\Phi_i$'s and compare the results obtained for various $\\Phi_i$'s with exact relativistic calculations. As accretion onto a black hole is {\\it necessarily} transonic to satisfy the inner boundary condition at the event horizon, unlike Bondi's (1952) original work, here we will concentrate only on transonic solutions, i.e., we will deal with accretion (wind) which is subsonic (supersonic) far away from the black hole and approaches (moves away from) the hole supersonically (subsonically) after crossing a sonic point (the location of which can be determined as a function of ${\\cal E}$ and $\\gamma$, see \\S 3) on its way towards (away from) the accretor. ",
        "conclusions": "In this paper we have solved a set of algebraic and differential equations governing various dynamical and thermodynamic behaviouars of Bondi (1954) type accretion in a full Schwarzschild metric as well as for motion under a number of `pseudo-Schwarzschild' potentials, to examine the suitability in application of those potentials in investigating spherically symmetric transonic accretion onto a non- rotating black hole. We have shown that though the potentials discussed here were originally proposed to mimic the relativistic effects manifested in the disc accretion, it is quite reasonable to use most of the potentials in studying various dynamical as well as thermodynamic quantities for spherical accretion. Also, we have shown that depending on the chosen regions of parameter space spanned by specific energy ${\\cal E}$ and adiabatic index $\\gamma$ of the flow, one potential may be important than others and we could identify which potential is the best approximation for FGRSF for what values of ${\\cal E}$ and $\\gamma$. We have restricted ourselves to the study of simple polytropic flows only. However, the validity of using all these $\\Phi_i$s discussed here can easily be examined for isothermal accretion and wind as well as for flows with other equations of state. Work in this direction is reported elsewhere (Sarkar \\& Das, 2001).\\\\ \\noindent It is observed that among all pseudo potentials, $\\Phi_1$ (potential proposed by Paczy\\'nski and Wiita (1980)) and $\\Phi_4$ (one of the potentials proposed by Artemova et. al. (1996)) are in general the best in the sense that they provide very reasonable approximation to the full general relativistic solution. While $\\Phi_1$ is the best approximation for ultra-relativistic flow, $\\Phi_4$ happens to be the best approximation as the flow tends to be fully non relativistic, i.e, $\\gamma$ tends to have the value $\\frac{5}{3}$. Also we see that there are certain cases for which one or more of the pseudo potentials may give the exact match with FGRSF for a particular value of ${\\cal E}$ or $\\gamma$ (for a fixed $r$) in finding some dynamical ($r_c, M$ etc.)  or thermodynamic (flow temperature $T$, for example) quantity.\\\\ \\noindent It is worth mentioning that as long as one is not interested in astrophysical processes extremely close (within $1-2~r_g$) to a black hole horizon, one may safely use the `pseudo' potentials discussed here to study spherically symmetric accretion on to a Schwarzschild black hole with the advantage that use of these potentials would simplify calculations by allowing one to use some basic features of flat geometry (additivity of energy or de-coupling of various energy components, i.e., thermal ($\\frac{a^2}{\\gamma-1}$) Kinetic ($\\frac{u^2}{2}$) or gravitational ($\\Phi$) etc.) which is not possible for calculations in a purely Schawarzschild metric. Also, one can study more complex many body problems such as accretion from an ensemble of companions or overall effeciency  of accretion onto an ensemble of black holes in a galaxy  or  for studying numerical hydrodynamic or magnetohydrodynamic flows around a black hole etc. as simply as can be done in a Newtonian framework, but with far better accuracy. However, one should be careful in using these potentials to study spherically symmetric accretion because of the fact that none of the potentials discussed here are `exact' in a sense that they are not directly derivable from the Einstein equations. These potentials could only be used to obtain more accurate correction terms over and above the pure Newtonian results and any `radically' new results obtained using these potentials should be cross-checked very carefully with the exact general relativistic theory.\\\\ \\noindent Although the theory of disc accretion has priority over spherical accretion because of the fact that accretion discs describe more realistic situations found in nature, it is not unreasonable to concentrate on spherical accretion because for certain cases, that may be quite useful and use of these potentials makes a complicated problem simpler to study. For example, for a supermassive black hole immersed in intergalactic space in such a way that matter falling on to it has negligible intrinsic angular momentum, the accretion (at least close to the hole) is quasi spherical and transonic spherical accretion might be a good approximation to mimic the situation. Same sort of approximation is valid when an accreting black hole is embedded in a number of donor stars (or star clusters)  where the angular momentum of the stars are randomly oriented in such a way that the vector sum of the intrinsic angular momentum carried by the accreting matter as a whole may be quite negligible, so as to make Bondi-type accretion a good approximation. In fact, a number of recent works (Coker \\& Markoff 2001 and references therein, Das, 1999, 2000, 2001a,b, Toropin, et. al, 1999, Kovalenko \\& Eremin 1998, Titarchuk, et. al. 1996, 1997, Wang \\& Sutherland 1997, Zampieri, et. al. 1996, Yim \\& Park 1995, Markovic 1995, Tsuribe, et. al. 1995, Kazhdan \\& Murzina 1994, Fortner 1993) still deal with spherical accretion to investigate some basic astrophysical processes the black holes and neutron stars. So we believe that work presented in this paper is relevant and will be useful in investigation of various aspects of accretion and wind around non-rotating and slowly rotating compact objects."
    },
    "0212/astro-ph0212282_arXiv.txt": {
        "abstract": "We present a detailed analysis of the dynamical properties of a simulated disk galaxy assembled hierarchically in the $\\Lambda$CDM cosmogony.  At $z=0$, two distinct dynamical components are easily identified solely on the basis of the orbital parameters of stars in the galaxy: a slowly rotating, centrally concentrated spheroid and a disk-like component largely supported by rotation. These components are also clearly recognized in the surface brightness profile of the galaxy, which can be very well approximated by the superposition of an $R^{1/4}$ spheroid and an exponential disk. However, neither does the dynamically-identified spheroid follow de Vaucouleurs' law nor is the disk purely exponential, a result which calls for caution when estimating the importance of the disk from traditional photometric decomposition techniques. The disk may be further decomposed into a thin, dynamically cold component with stars on nearly circular orbits and a hotter, thicker component with orbital parameters transitional between the thin disk and the spheroid. Supporting evidence for the presence of distinct thick and thin disk components is found, as in the Milky Way, in the double-exponential vertical structure of the disk and in abrupt changes in the vertical velocity distribution as a function of age.  The dynamical origin of these components offers intriguing clues to the assembly of spheroids and disks in the Milky Way and other spirals. The spheroid is old, and has essentially no stars younger than the time elapsed since the last major accretion event; $\\sim 8$ Gyr ago for the system we consider here. The majority of thin disk stars, on the other hand, form after the merging activity is over, although a significant fraction ($\\sim 15\\%$) of thin-disk stars are old enough to predate the last major merger event. This unexpected population of old disk stars consists mainly of the tidal debris of satellites whose orbital plane was coincident with the disk and whose orbits were circularized by dynamical friction prior to full disruption. More than half of the stars in the thick disk share this origin, part of a trend that becomes more pronounced with age: nine out of ten stars presently in the old ($\\tau \\gsim \\, 10$ Gyr) disk component were actually brought into the disk by satellites. By contrast, only one in two stars belonging to the old spheroid are tidal debris; the rest may be traced to a major merger event that dispersed the luminous progenitor at $z\\sim 1.5$ and seeded the formation of the spheroid. Our results highlight the role of satellite accretion events in shaping the disk---as well as the spheroidal---component and reveal some of the clues to the assembly process of a galaxy preserved in the detailed dynamics of old stellar populations. ",
        "introduction": "\\label{sec:intro}  Some of the observed properties of galactic disks seem at odds with many of the ``natural'' trends expected in hierarchically clustering models and present a significant challenge to the current paradigm of structure formation on small scales (see, e.g., Sellwood \\& Kosowsky 2001, Navarro \\& Steinmetz 2000, Moore 2001).  Qualitatively, the main difficulty lies in reconciling the early collapse and eventful merging history characteristic of the buildup of galactic dark matter halos with the many dynamical clues which point to a smooth assembly of the luminous component of galactic disks. This difficulty has made it necessary to postulate a substantial role for complex astrophysical processes in order to overcome some of these trends and to bring models of hierarchical assembly into agreement with observation.  The prime concern regards the fragility of centrifugally supported stellar disks to rapid fluctuations in the gravitational potential such as those stirred by mergers and satellite accretion events (Toth \\& Ostriker 1992, Quinn, Hernquist \\& Fullagar 1993, Velazquez \\& White 1999).  As first computed in detail by T\\'oth \\& Ostriker (1992), the constraints on merger events undergone by a thin disk such as that of the Milky Way are very strict indeed. These authors find that less than $10\\%$ of the disk mass within the solar circle could have been accreted in the past $5$ Gyr in the form of clumps, limiting severely the role of merging in the recent mass accretion history of the Milky Way.  Although these numbers might be relaxed somewhat by taking into account the coherent response of a self-gravitating disk to the accretion event (Huang \\& Carlberg 1997, Walker, Mihos \\& Hernquist 1996, Vel\\'azquez \\& White 1999), there is broad consensus that the presence of a dominant, thin, cold stellar disk points convincingly to a history of mass accretion where major mergers have played a minor role. Indeed, since the fragile circular orbits of disk stars are easily perturbed during mergers, the age of the oldest disk stars is often used to estimate the epoch of the last major merger. In the case of the Milky Way, stars as old as $\\sim 14$ Gyr are found in the solar neighborhood, suggesting a protracted (and perhaps episodic) history of star formation in the disk (Wyse 2000, Rocha-Pinto et al 2000, Liu \\& Chaboyer 2000) and, consequently, a paucity of merger events at odds with typical merging histories of dark halos in cosmogonies such as the $\\Lambda$CDM model.   The need for ordered and smooth settling of gas into thin galactic disks leads to a few basic predictions of hierarchical scenarios that can be contrasted directly with observation. For example, because angular momentum results from torques operating before collapse, it correlates directly with the radius and time of turnaround (Navarro \\& Steinmetz 1997). Thus, material that turns around later is expected to have higher angular momentum, with two major implications: (i) disks are expected to be physically smaller in the past, compared with systems of similar rotation speed identified at present, and (ii) disks form from the inside out, as higher angular momentum material, which collapses later, should settle preferentially in the outer regions of the disk (Mo, Mao \\& White 1998).  Although overall these expectations are not grossly inconsistent with current data, a number of well established observations are worryingly difficult to accommodate within this general scenario.  One of these worries concerns the origin of the complex vertical structure of galactic disks. Current datasets suggest that most (if not all) galactic disks are built out of two major dynamical components, referred to generally as the ``thick'' and ``thin'' disks (Dalcanton \\& Bernstein 2002 and references therein). Although varying in importance from galaxy to galaxy, the thick disk appears to be old ($\\gsim 10$-$12$ Gyr in the case of the Milky Way, Gilmore, Wyse \\& Jones 1995), poor in metals (Prochaska et al 2000), and of similar radial extent as the thin disk component (Wyse 2000).  These trends have led to the popular conjecture that the Milky Way's thick disk has its origin in an early thin disk of velocity and size comparable to today's but ``thickened'' by the accretion of a satellite $\\sim 10$-$12$ Gyr ago (Gilmore, Wyse \\& Norris 2002). Appealing as this idea may be from a dynamical standpoint, it is rather unattractive in a hierarchical scenario, where the mere existence of such large, fast-rotating disks at $z\\sim 2$ (corresponding to roughly $\\sim 12$ Gyr ago in the $\\Lambda$CDM scenario) counters the ``natural'' prediction of the model.  A related difficulty concerns the presence of old disk stars in the vicinity of the Sun, which might be taken to imply that even at high redshift the thin disk already extended as far as the solar circle and has remained essentially undisturbed dynamically since. Such a long period of quiet dynamical evolution is difficult to reconcile with the active merging expected at early times in the $\\Lambda$CDM cosmogony, and perhaps even with the presence of the Galaxy's spheroidal component, which is typically ascribed to a relatively recent major merger event.  The discussion above illustrates the still unsettled account of the origin of the Galactic disk(s) in the hierarchical formation paradigm, and draws notice to the fundamental role that accretion events, be they mergers or satellite disruptions, have played in determining the structure of galactic disks. In this paper---the second in a series analyzing the dynamical properties of galaxies simulated in the ``concordance'' $\\Lambda$CDM paradigm---we examine these issues using a numerical simulation with high numerical resolution. The simulation is the same as described by Abadi et al (2002) in Paper I of this series; that paper confronts the global dynamical and photometric properties of the simulated galaxy with normal spirals, while the present one concentrates on the multi-component nature of the stellar disk and on its dynamical origin.  The paper is organized as follows. In section~\\ref{sec:numexp} we briefly discuss, for completeness, numerical details of the simulation; \\S~\\ref{sec:results} discusses the identification of distinct dynamical components in the simulated galaxy, as well as the implications of these results in the context of the formation of the thick and thin disks of the Milky Way. Finally, \\S~\\ref{sec:conc} summarizes our main conclusions. ",
        "conclusions": "\\label{sec:conc}  We present a detailed analysis of the dynamical components of a galaxy simulated in the ``concordance'' $\\Lambda$CDM cosmogony. The galaxy forms in a dark matter halo chosen so that mergers and accretion events are unimportant dynamically after $z\\sim 1$. Star formation and feedback parameters are such that the star formation history of the galaxy is largely driven by the rate at which gas cools and collapses within dark halos: this is a conservative choice where feedback effects are relatively unimportant. The shortcomings of this assumption concerning the global photometric and structural properties of the simulated galaxy are discussed in Paper I. We focus here on the multi-component nature of the stellar disk and on their dynamical origin.  Our main results can be summarized as follows.  \\begin{itemize}  \\item{At $z=0$, two separate stellar components are easily distinguishable solely on the basis of the orbital parameters of stars in the galaxy: a slowly rotating, centrally concentrated spheroid and an extended disk-like component largely supported by rotation. }  \\item{These components are also recognized in the surface brightness profile of the galaxy, which can be very well approximated by the superposition of an $R^{1/4}$ spheroid and an exponential disk, in agreement with observations of early-type spirals.}  \\item{ Neither does the dynamically identified spheroid follow closely de Vaucouleurs' law nor is the disk purely exponential, a result which calls for caution in the interpretation of estimates of the dynamical importance of the disk and spheroid in traditional photometric decomposition techniques.}  \\item{The spheroid is old, and has essentially no stars younger than the last major accretion episode; $\\sim 8$ Gyr ago for the system we consider here. The majority of thin disk stars, on the other hand, form after the merging activity is over and have a mean age of $\\sim 5$ Gyr.}  \\item{The disk may be further decomposed into two well defined subcomponents: a thin, dynamically cold disk of stars on nearly circular orbits and a thicker disk with orbital parameters transitional between the thin disk and the spheroid.  Supporting evidence for the true presence of these two distinct components is found, as in the Milky Way, in the double-exponential vertical structure of the disk and in abrupt changes in the vertical velocity distribution as a function of age.}  \\item{The bulk ($\\sim 60\\%$) of the thick disk consists of the tidal debris of satellites whose orbital plane was coincident with the disk and whose orbits were circularized by dynamical friction prior to full disruption. This trend becomes more pronounced with age; $\\sim 90\\%$ of stars older than $10$ Gyr and presently in the thick disk component were brought into the disk by satellites.}  \\item{ A significant fraction ($\\sim 15\\%$) of thin-disk stars are old enough ($\\gsim \\, 10$ Gyr) to predate the last major accretion event. The bulk of this unexpected population of old stars on nearly circular orbits share a common origin with the old thick disk: they are the remains of the cores of disrupted satellites. Interestingly, only one in two stars belonging to the old spheroid are tidal debris; the rest may be traced to a major merger event that disperses the luminous progenitor at $z\\sim 1.5$ and seeds the formation of the spheroid.}  \\end{itemize}   Our results offer clues to understanding a number of observational trends that challenge the standard hierarchical disk assembly process. In particular, the presence of an old disk component made up primarily of tidal debris from disrupted satellites helps to explain: (i) the presence of a significant number of old stars on circular orbits in the outskirts of galaxies like the Milky Way, and (ii) why the specific angular momentum and radial extent of the thick and thin disk components are comparable in spite of the significantly different ages of their individual stars. It also offers a ``natural'' explanation for the clear dynamical and evolutionary distinction between the thin and thick disk components: the thick disk is mostly tidal debris from disrupted satellites whereas the young thin disk consists mostly of stars formed ``in situ'' after the merging activity abates.  These findings highlight the role of satellite accretion events in shaping the disk---as well as the spheroidal---components of a galaxy and reveal some of the clues to the assembly process of the galaxy preserved within the detailed dynamics of old stellar populations. We emphasize that these conclusions are based on the detailed analysis of a single simulation and that further work is needed to put these results on a firmer basis as well as to establish their true applicability to the origin of the dynamical components of the Milky Way. We plan to explore these issues in detail in future papers of this series."
    },
    "0212/astro-ph0212557_arXiv.txt": {
        "abstract": "The $\\VVM$ of the cosmological X-ray flashes  detected by Wide Field Cameras on {\\it BeppoSAX} is calculated theoretically in  a simple jet model. The total emission energy from the jet is assumed to be constant. We find that if the jet opening half-angle is smaller than 0.03 radian, the theoretical  $\\VVM$ for fixed opening half-angle is less than $\\sim0.4$, which is consistent with the recently reported observational value of $0.27\\pm 0.16$ at the 1 $\\sigma$ level. This suggests  that the off-axis GRB jet with the  small opening half-angle at the cosmological distance can be identified as the cosmological X-ray flash. ",
        "introduction": "The X-ray flash (XRF) is a class of X-ray transients (Heise et al. 2001, see also Barraud et al. 2003). Some properties of XRFs, such as the observed event rate, the duration and the isotropic distribution, are similar to that of Gamma-Ray Bursts (GRBs), while the spectral hardness of XRFs characterized by the peak flux ratio, the fluence ratio and the photon index is softer than that of GRBs. This class represents a large portion of the whole GRB population. Recently, the observational value of $\\VVM$, which is the measure of the homogeneity of spatial distribution (Schmidt, Higdon, \\& Hueter 1988; see also Chang \\& Yi 2001; Kim, Chang, \\& Yi 2001), has been updated from $0.56\\pm0.12$ (J.~Heise 2000, talk given in 2nd workshop Gamma-Ray Bursts in the Afterglow Era) to $0.27\\pm0.16$ (J.~Heise 2002, talk given in 3rd workshop Gamma-Ray Bursts in the Afterglow Era). The updated value of $\\VVM$ suggests that XRFs take place at a cosmological distance.  Various models accounting for the nature of the XRFs have been proposed (Yamazaki, Ioka, \\& Nakamura 2002; Heise et al. 2001; Dermer, Chiang, \\& B\\\"{o}ttcher 1999; Huang, Dai, \\& Lu 2002; M\\'{e}sz\\'{a}ros et al. 2002; Mochkovitch et al. 2003; Daigne \\& Mochkovitch 2003). Heise et al. (2001) proposed that XRFs could be {\\it GRBs at high redshift}. The redshifts of  XRF011030 and XRF020427 have an upper limit of $z\\lesssim 3.5$ (Bloom et al. 2003; J.~Heise 2002, talk given in 3rd workshop Gamma-Ray Bursts in the Afterglow Era), which suggests that XRFs take place at not so high redshift but the same as that of GRBs. {\\it The photosphere-dominated fireball model} may account the nature of the XRFs with peak energy $E_p$ more than $\\sim20$ keV (M\\'{e}sz\\'{a}ros et al. 2002, Ramirez-Ruiz \\& Lloyd-Ronning 2002). However, further considerations are needed to explain the event with $E_p\\sim $ of approximately a few keV, such as XRF020427 (L.~Amati 2002, talk given in 3rd workshop Gamma-Ray Bursts in the Afterglow Era), XRF020903 and XRF010213 (N.~Kawai 2002, talk given in 3rd workshop Gamma-Ray Bursts in the Afterglow Era). The models with small Lorentz factors, such as {\\it the dirty fireball model} (Dermer, Chang, \\& B\\\"{o}ttcher 1999; Huang, Dai, \\& Lu 2002) or {\\it the structured-jet model} (Rossi, Lazzati, \\& Rees 2001; Woosley et al. 2002; Zhang \\& M\\'{e}sz\\'{a}ros 2002a) also have possibilities to explain the properties of the XRF, with the implicit assumption that the XRF arises not from internal shocks (Zhang \\& M\\'{e}sz\\'{a}ros 2002b). The models for internal shocks with {\\it small contrast of high Lorents factors} might be the origin of XRFs (Mochkovitch et al. 2003; Daigne \\& Mochkovitch 2003).  For the other possibility, we have studied {\\it the off-axis jet model} and proposed that if we observe the GRB jet with a large viewing angle, it looks like an XRF (Yamazaki, Ioka, \\& Nakamura 2002, Yamazaki, Ioka, \\& Nakamura 2003). In Yamazaki, Ioka, \\& Nakamura (2002)  the value of the jet opening half-angle was  adopted as  $\\Delta\\theta=0.1$. In this  model  the distance to the farthest XRF ever detected is about 2 Gpc ($z\\sim0.4$) so that the cosmological effect is small and $\\VVM\\sim 0.5$. Recent observations suggest that GRBs with relatively small opening angle exist, while the distribution of $\\Delta\\theta$ is not yet clear (Panaitescu \\& Kumar 2002). If we assume the total emission energy to be constant as in the previous paper, the intrinsic luminosity is larger for the smaller opening half-angle. Such GRBs at the cosmological distance observed from off-axis viewing angle may be seen as XRFs and $\\VVM$ is expected to be  smaller than 0.5.   In this paper, we will show that our off-axis model has a possibility of accounting for the observational value of $\\VVM$ if we change some of the model parameters in the previous paper (Yamazaki, Ioka, \\& Nakamura 2002). This paper is organized as follows. In \\S~\\ref{sec:model}, we describe a simple jet model including the effect of cosmological expansion. We assume the uniform jet with sharp edges. Although one may consider the structured jet motivated by the simulation of the collapsar model, we cannot conclude, both observationally and theoretically, which model is preferable. The $\\VVM$ for the XRFs detected by the Wide Field Cameras (WFCs) on {\\it BeppoSAX} is calculated in \\S~\\ref{sec:vvm}. Section~\\ref{sec:dis} is devoted to a discussion. Throughout this paper, we adopt the flat universe with $\\Omega_M=0.3$, $\\Omega_\\Lambda=0.7$ and $h=0.65$. ",
        "conclusions": "\\label{sec:dis} We have calculated $\\VVM_{\\Delta\\theta}$ for the emission from a simple jet model, and shown that when the jet opening half-angle $\\Delta\\theta$ is smaller than about 0.03, $\\VVM_{\\Delta\\theta}$ for the XRFs detected by WFCs on {\\it BeppoSAX} is smaller than 0.4. The value of $\\Delta\\theta\\sim0.03$ has been obtained from the fitting of the afterglow light curve (Panaitescu \\& Kumar 2002; Frail et al. 2001). Such a narrow jet which is inferred by afterglow observations is rare. However, following Eq.~(1) of Frail et al. (2001), the jet break time is given by $t_j\\sim 13(\\Delta\\theta/0.01)^{8/3}$~minutes and so it requires fast localization to observe the jet break for a narrow jet. Therefore, at present, the small number of GRBs with small $\\Delta\\theta$ may come from the observational selection effect. In the context of this scenario, we might be able to account for the fact that afterglows of XRFs have been rarely observed since the afterglow at a fixed time gets dimmer for an earlier break time. Furthermore, some ``dark GRBs'' might be such a small opening angle jet observed with an on-axis viewing angle for the same reason.   We briefly comment on how the results obtained in this paper will depend on the Lorentz factor of the jet $\\gamma$. We see that when $\\gamma$ becomes large, $\\VVM_{\\Delta\\theta}$ becomes small. For example, when we fix $\\gamma=200$, we obtain $\\VVM_{\\Delta\\theta=0.03}=0.39$ and $\\VVM_{\\Delta\\theta=0.1}=0.43$. This implies that the limitation on $\\Delta\\theta$ can be relaxed.  We can estimate the typical observed photon energy as $h\\nu_{obs}\\sim (1+z)^{-1}\\delta\\nu'_0,$ where $\\delta^{-1}=\\gamma[1-\\beta\\cos\\tilde{\\theta}]$ and $\\tilde{\\theta}={\\rm max}\\{0,\\,\\theta_v-\\Delta\\theta\\}$ (Yamazaki, Ioka, \\& Nakamura 2002). Since $\\gamma\\gg 1$ and $\\theta_v,\\,\\Delta\\theta\\ll 1$, we obtain \\begin{equation} h\\nu_{obs}\\sim \\f{2\\gamma\\nu'_0}{(1+z) [1+(\\gamma\\tilde{\\theta})^2]}. \\end{equation} In \\S~\\ref{sec:vvm}, we have shown that for fixed $\\Delta\\theta\\lesssim0.03$, the typical value of $\\theta_v$ is $\\sim\\Delta\\theta+0.02$. Therefore, for the adopted parameters $\\gamma\\nu'_0=200$keV and the typical redshift $z=z_p\\sim 1.5$, one can derive $h\\nu_{obs}\\sim 30$ keV, which is the typical observed peak energy of the XRFs (Kippen et al. 2002). We can propose from our argument that the emissions from the jets with a small opening half-angle such as $\\Delta\\theta\\lesssim0.03$ are observed as XRFs when they are seen from off-axis viewing angle. If one can detect the afterglow of the XRF, which has the maximum flux at about several hours after the XRF, the fitting of light curve may give us the key information about the jet opening angle (Granot et al. 2002). Therefore, our theoretical model can be tested by the near-future observations.  We can estimate the observed event rate of the XRF for fixed $\\Delta\\theta$ as, \\begin{equation} R_{{\\rm XRF},\\,\\Delta\\theta}=\\f{1}{4\\pi}\\int W(\\theta_v)\\,d\\theta_v. \\end{equation} In order to calculate $R_{{\\rm XRF},\\,\\Delta\\theta}$, we consider the proportionality constant ${\\cal R}=n(z)/n_{SF}(z)$. One can write approximately as ${\\cal R}=r_{>8M_\\sun}k$, where $r_{>8M_\\sun}$ is the number of stars with masses $M>8M_\\sun$ per unit mass. Since we assume all stars with masses $M>8M_\\sun$ explode as core-collapse SNe, $k$ represents the ratio of the number of XRF sources to that of core-collapse SNe. We adopt the value $k=1\\times10^{-3}$, which is derived by the result of Porciani \\& Madau (2001) combined with the effect of the solid angle factor $(\\Delta\\theta)^2/2$. Using a Salpeter initial mass function $\\phi(M)$, we obtain $r_{>8M_\\sun}=(\\int_{8M_\\sun}^{125M_\\sun}\\phi(M)\\,dM)/ (\\int_0^{125M_\\sun}M\\phi(M)\\,dM) =1.2\\times10^{-2}M_\\sun^{-1}$. Then, in the case of $\\Delta\\theta=0.03$, we derive $R_{{\\rm XRF},\\,\\Delta\\theta}=1\\times 10^2 \\ {\\rm events} \\ {\\rm yr}^{-1} ({\\cal R}/1\\times10^{-5}M_\\sun^{-1})$, which is comparable to the observed event rate of the XRF (Heise et al. 2001). Note that the value of $R_{{\\rm XRF},\\,\\Delta\\theta}$ remains unchanged within a factor of 2 when we vary $\\Delta\\theta$ from 0.01 to 0.07.  When the jet opening half-angle $\\Delta\\theta$ has a distribution $f_{\\Delta \\theta}$, we integrate $\\VVM_{\\Delta\\theta}$ and $R_{{\\rm XRF},\\,\\Delta\\theta}$ over the distribution of $\\Delta\\theta$ as \\begin{equation} \\VVM= \\f{\\int d(\\Delta\\theta)\\, f_{\\Delta\\theta}\\,R_{{\\rm XRF},\\,\\Delta\\theta}\\, \\VVM_{\\Delta\\theta}} {\\int d(\\Delta\\theta)\\,f_{\\Delta\\theta}\\, R_{{\\rm XRF},\\,\\Delta\\theta}} \\ \\ , \\end{equation} \\begin{equation} R_{{\\rm XRF}}=\\f{\\int d(\\Delta\\theta)\\,f_{\\Delta\\theta}\\, R_{{\\rm XRF},\\,\\Delta\\theta}} {\\int d(\\Delta\\theta)\\,f_{\\Delta\\theta}} \\ \\ , \\end{equation} respectively. We assume a power-low distribution as $f_{\\Delta\\theta}\\propto(\\Delta\\theta)^{-q}$. When we adopt $q=4.54$ (Frail et al. 2001), and integrate over $\\Delta\\theta$ from 0.01 to 0.2 rad, we find $\\VVM=0.36$ and $R_{{\\rm XRF}}=1\\times10^2$~events~yr$^{-1}$. These values mainly depend on the lower bound of the integration. For example, we obtain $\\VVM=0.43$ and $R_{{\\rm XRF}}=3$~events~yr$^{-1}$ if the integration is done over $\\Delta\\theta$ from 0.03 to 0.2 rad (Note that we may let the value of $R_{\\rm XRF}$ be consistent with observed value by adjusting ${\\cal R}$). Since the statistics of the observational data will increase in the near future owing to instruments such as {\\it HETE-2} and {\\it Swift}, we will be able to say more than above discussion, including more accurate functional form of $f_{\\Delta\\theta}$ than that we have considered above, as well as the relation to the GRB event rate."
    },
    "0212/astro-ph0212411_arXiv.txt": {
        "abstract": "s{As part of a research of elusive AGN in a well defined sample of infrared selected galaxies (using {\\em Beppo}SAX, Chandra, {\\em XMM-Newton} and INTEGRAL data) we have observed with \\sax the merging starburst system Arp~299. The broad-band (0.1--40 keV) coverage of this observation clearly reveals, for the first time in this system, the presence of a deeply buried AGN having an intrinsic luminosity of $L_{0.5-100\\:\\kev} \\simeq 1.9 \\times 10^{43}\\:$\\lum .} ",
        "introduction": "Studies of active objects at IR and X-ray wavelengths indicate that star-formation and AGN activity may be related (Fadda et al. 2002). The trigger mechanism for both phenomena could be the interaction or the merging of gas-rich galaxies. This generates fast compression of the available gas in the inner galactic regions, causing both the onset of a major starburst and the fueling of a central black hole raising the AGN activity.  However the concomitant AGN and starburst activity is expected to happen in a high-density medium ($N_H \\geq 10^{23-24}\\:$\\nh), characterized by high dust extinction of the UV-optical flux and strong photoelectric absorption of the soft X-rays (e.g. Fabian et al. 1998). Thus the study of these active phases in galaxies becomes very difficult; optical and even mid-/far-IR spectroscopy may not be sufficient to disentangle starburst activity from AGN activity, which is actually best probed in the hard ($E > 6\\:$keV, in order to sample also the Fe K$\\alpha$ line) X-ray energy band.  To search for hidden AGNs and to shed light on the starburst-AGN connection and its occurrence we have started a systematic and objective investigation in hard ($E >6\\:$keV) X-rays of {\\it a flux-limited sample of IRAS galaxies}. The sample consists of 28 galaxies selected from the {\\it IRAS Cataloged Galaxies and Quasars} (see http://irsa.ipac.caltech.edu/) as having $f_{60\\mum} > 50\\:$Jy or $f_{25\\mum} > 10\\:$Jy. We stress here that no other selection criteria (e.g. established presence of an AGN, luminosities, IR colours, etc.) have been applied to the sample definition. ",
        "conclusions": ""
    },
    "0212/astro-ph0212405_arXiv.txt": {
        "abstract": "The morphological dependence of the luminosity function is studied using a sample containing approximately 1500 bright galaxies classified into  Hubble types by visual inspections for a homogeneous sample obtained from the Sloan Digital Sky Survey (SDSS) northern equatorial stripes. Early-type galaxies are shown to have a characteristic magnitude by 0.45 mag brighter than spiral galaxies in the $r^{\\ast}$ band, consistent with the `universal characteristic luminosity' in the $B$ band. The shape of the luminosity function differs rather little among different morphological types: we do not see any symptoms of the sharp decline in the faint end for the luminosity function for early-type galaxies at least 2 mag fainter than the characteristic magnitude, although the faint end behaviour shows a slight decline ($\\alpha\\lsim -1$) compared with the total sample. We also show that a rather flat faint end slope for early-type galaxies is not due to an increasing mixture of the dwarf galaxies which have softer cores. This means that there are numerous faint early-type galaxies with highly concentrated cores. ",
        "introduction": "The origin of morphology of galaxies is a long-standing issue, which could provide a key to discerning among models of the formation of galaxies. How galaxy morphology changes as a function of the lookback time is perhaps the prime approach to this problem, and the knowledge of the morphological dependence of the local luminosity function at zero redshift is the baseline. One specific example of the issues is whether the luminosity function of elliptical galaxies obeys the Schechter-type function with a rather flat faint end (e.g., Marzke et al. 1994; Kochanek et al. 2001), or the Gaussian function as inferred by Binggeli, Sandage \\& Tammann (1988) and more recently by Bernardi et al. (2002a). If the latter is correct, one would envisage galaxy morphology as a bulge-luminosity sequence (Dressler \\& Sandage 1983; Meisels and Ostriker 1984), which in turn reveals a clue about the formation of elliptical galaxies and bulges.  There are also a number of uses of the morphology-dependent luminosity function (MDLF). We mention only one example: the frequency of gravitational lensing of quasar images is approximately proportional to the luminosity density of early-type galaxies rather than that of all galaxies (Fukugita \\& Turner 1991). The uncertainty in the MDLF is the largest source of error in predicting the frequency of gravitational lenses, and thus in inferring the cosmological constant from such analyses.  The understanding of the MDLF is significantly poorer than that of the luminosity function for galaxies in general, which has undergone substantial progress in the latest years (Folkes et al. 1999; Blanton et al. 2001) by virtue of large galaxy samples. The traditional way to obtain MDLF is to use morphological classification based on visual inspections of the images (Binggeli et al. 1988; Loveday et al. 1992; Marzke et al. 1994; 1998; Kochanek et al. 2001). Some modern  studies attempt to use spectroscopic features to classify galaxies into morphological types (Bromley et al. 1998; % Folkes et al. 1999), which makes it possible to analyze large samples. Although a general correlation is known between spectroscopic and Hubble morphologies, the samples derived from the two methods are considerably different. In particular, the classification using spectroscopic features or colours is sensitive to small star formation activities now or in the near past in early-type galaxies, whilst Hubble morphology is insensitive to this process. The problem of automated classification always lies in the difficulty in finding quantitative measures that strongly correlate with the Hubble sequence based on visual inspections. In this paper we derive the MDLF based on visual classifications using a homogeneous bright galaxy sample from Sloan Digital Sky Survey (SDSS; York et al. 2000). The sample we use in this paper is small, but it is based on a homogeneous morphological classification with accurate photometry.  The SDSS conducts both photometric (Gunn et al. 1998; Hogg et al. 2001; Pier et al. 2002) and spectroscopic surveys, and is producing a homogeneous data set, which is suitable to studies of galaxy statistics. The initial survey observations were made in the northern and southern equatorial stripes, and produced a galaxy catalogue to $r^{\\ast}=22.5$ mag in five colour bands (Fukugita et al. 1996) with a photometric calibration using a new standard star network observed at USNO (Smith et al. 2002). Spectroscopic follow-up is made to 17.8 mag with accurately defined criteria for target selection (Strauss et al. 2002).  Our study is limited to bright galaxies with $r^{\\ast}\\leq 15.9$ mag after Galactic extinction correction, since visual classifications sometimes cannot be made confidently beyond this magnitude with the SDSS imaging data. We have classified all galaxies satisfying this magnitude criterion in the northern equatorial stripe. The total number of galaxies in our sample is 1875, of which 1600 have spectroscopic information.  The dominant part of the data we used are already published as an {\\it Early Data Release} (EDR) (Stoughton et al. 2002). Our present work uses primarily the EDR but supplemented by observations which are not included in EDR to make the sample as complete as possible. Photometry of galaxies in this region is discussed in a galaxy number count paper of Yasuda et al. (2001), and the luminosity function is derived by Blanton et al. (2001), which also discuss spectroscopic details. ",
        "conclusions": "Our sample is small and we may not be able to extract quantitatively robust parameters, yet we obtain a number of useful conclusions. The most important feature with our analysis is that we have used a homogeneous photometric catalogue with sharply defined selection criteria and a homogeneously morphologically-classified sample based on Hubble morphology of galaxies, rather than a sample classified by indicators using spectroscopic features or colours, which are sensitive to small star formation activities in the present or the near past.  The first conclusion we have obtained is that the shape of the MDLF does not depend too strongly on the Hubble types. The characteristic luminosity of elliptical and S0 galaxies is brighter than that of spiral galaxies in the $r^{\\ast}$ band. The amount of the difference in brightness is consistent with universal characteristic luminosity in the $B$ band, which was found by Tammann et al. (1979). The MDLF of early-type galaxies somewhat declines in the faint end, but does not exhibit a sharp decline, and this is not due to an increasing mixture of dwarf galaxies at least in the magnitude range we are concerned with. The conclusion is unchanged if we use the concentration index as a classifier of early-type galaxies. This indicates that there are many intrinsically faint elliptical galaxies, whose luminosities are fainter than those of bulges in spiral galaxies. The existence of numerous early-type galaxies with a hard core at small luminosities indicates that morphology is unlikely to be a sequence of the bulge luminosity as advocated by Dressler \\& Sandage (1983), and by Meisels and Ostriker (1984). Our conclusion also justifies the calculation of the strong gravitational lensing frequency of quasars using the standard Schechter function without introducing a cutoff in the luminosity function, which would affect the frequency of sub-arcsecond lensing.  \\vspace{10pt}   Funding for the creation and distribution of the SDSS Archive has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Aeronautics and Space Administration, the National Science Foundation, the U.S. Department of Energy, the Japanese Monbukagakusho, and the Max Planck Society. The SDSS Web site is http://www.sdss.org/. The SDSS is managed by the Astrophysical Research Consortium (ARC) for the Participating Institutions. The Participating Institutions are The University of Chicago, Fermilab, the Institute for Advanced Study, the Japan Participation Group, The Johns Hopkins University, Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, University of Pittsburgh, Princeton University, the United States Naval Observatory, and the University of Washington. We would like to thank Sadanori Okamura for useful comments. MF is supported in part by the Grant in Aid of the Japanese Ministry of Education.  \\clearpage"
    },
    "0212/astro-ph0212319_arXiv.txt": {
        "abstract": "{First results from a deep XMM-Newton observation of a field in the Large Magellanic Cloud (LMC) near the northern rim of the supergiant shell LMC\\,4 are presented. Spectral and temporal analyses of a sample of selected X-ray sources yielded two new candidates for supernova remnants, a supersoft X-ray source and a likely high mass X-ray binary (HMXB) pulsar. From the fourteen brightest sources up to ten are active galactic nuclei in the background of the galaxy which can be used as probes for the interstellar medium in the LMC. From the three previously known HMXBs the Be/X-ray binary \\exo\\ was the brightest source in the field, allowing a more detailed analysis of its X-ray spectrum and pulse profile. During the pulse \\exo\\ shows eclipses of the X-ray emitting areas with increased photo-electric absorption before and after the eclipse. The detection of X-ray pulsations with a period of 69.2 s is confirmed for \\rxbe\\ and a possible period of 272 s is discovered from \\xmmbex. The results are discussed with respect to individual sources as well as in the view of source population studies in the vicinity of the supergiant shell LMC\\,4. ",
        "introduction": "Due to their proximity the Large and Small Magellanic Cloud (LMC, SMC) were always subject to X-ray surveys for satellites carrying imaging instruments. The most sensitive and most complete survey is available from ROSAT data obtained between 1990 and 1998 in the energy range 0.1--2.4 keV. In particular pointed observations with the PSPC detector with its large field of view covered in total $\\sim$59 and $\\sim$18 square degrees of the LMC and SMC regions on the sky \\citep{1999A&AS..139..277H,2000A&AS..142...41H}. The ROSAT PSPC and HRI observations revealed about 1000 and 750 X-ray sources in the direction of the LMC and SMC, respectively \\citep[see also][]{2000A&AS..143..391S,2000A&AS..147...75S}.  The large number of X-ray sources in the Magellanic Cloud fields allows population studies on rich samples of various kinds of source types powered by different mechanisms. An unusually high concentration of Be/X-ray binary systems was found in the SMC \\citep{2000A&A...359..573H}. Detectors sensitive to higher energies on board ASCA and XTE detected pulsations from many new Be/X-ray transients during outburst \\citep{2002Yokogawa}. The high sensitivity of the XMM-Newton instruments promises the detection of pulsations to much lower X-ray fluxes \\citep{2001A&A...369L..29S} usually observed from Be/X-ray binaries in their low state.  To study the source population in the LMC to flux limits below \\oergcm{-14} a deep $\\sim$60 ks XMM-Newton observation was performed as part of the telescope scientist guaranteed time. The selected field, aimed at RA = 05$^{\\rm h}$31$^{\\rm m}$20$^{\\rm s}$ and Dec = --65\\degr57\\arcmin38\\arcsec, was chosen because of its location on the rim of the supergiant shell (SGS) LMC\\,4 \\citep{1980MNRAS.192..365M} and the simultaneous coverage of three known high mass X-ray binaries (HMXBs). These are : 1) The Be/X-ray binary pulsar \\exo\\ discovered in outburst during EXOSAT observations of the LMC\\,X-4 region in 1983 \\citep{1985SSRv...40..379P}. X-ray pulsations with a pulse period of 13.7 s were detected for the first time in ROSAT data by \\citet{1996rftu.proc..131D}. 2) An outburst observed from \\rxbe\\ by ROSAT and the detection of 69 s X-ray pulsations led to the suggestion by \\citet{1997A&A...318..490H} that \\rxbe\\ is also a Be/X-ray binary. Optical spectroscopy by \\citet{2002A&A...385..517N} confirmed this. 3) The highly variable source \\rxob\\ proposed as first HMXB in the LMC powered by accretion from the wind of an OB supergiant companion \\citep{1995A&A...303L..49H} which is supported by optical spectroscopy \\citep{2002A&A...385..517N}.  Active galactic nuclei (AGN) behind nearby galaxies are of interest for two main reasons. They can be used to define a precise reference coordinate system \\citep[e.g.][]{2000AJ....120..845A} and provide line of sight probes to study the interstellar medium in the galaxies \\citep[e.g.][]{2001A&A...371..816K,2001A&A...365L.208H}. The new X-ray observatories Chandra and XMM-Newton will largely increase the number of background AGN as first results from LMC observations demonstrate \\citep{2001A&A...365L.208H,2002ApJ...569L..15D}.  In this article first results from the deep XMM-Newton observation are presented. The paper concentrates on spectral and temporal analyses of a selected sample of X-ray sources which includes the three known HMXBs and new candidates for supernova remnants (SNRs), a supersoft source (SSS), a likely Be/X-ray binary pulsar and ten candidates for background AGN.  \\begin{figure*} \\resizebox{\\hsize}{!}{\\includegraphics[clip=]{rgb_final.ps}} \\caption{EPIC-pn RGB color image produced from data in three energy bands (red: 0.3--1.0 keV, green: 1.0--2.0 keV and blue: 2.0--7.5 keV). Each individual image was smoothed with an adaptive intensity filter. Instrumental background from an observation with closed filter wheel (from revolution 59) was subtracted before vignetting and exposure correction. The XMM-Newton pointing was aimed at RA = 05$^{\\rm h}$31$^{\\rm m}$20$^{\\rm s}$ and Dec = --65\\degr57\\arcmin38\\arcsec\\ and covers about 13\\arcmin\\ in radius. North is to the top and West to the right. The brightest source, located in the South, is the Be/X-ray binary system \\exo.} \\label{epic-ima} \\end{figure*} ",
        "conclusions": "The deep XMM-Newton observation of a northern field in the LMC allows to uniquely classify the detected X-ray sources to flux levels of $\\sim$\\oergcm{-14} from their X-ray properties alone. X-ray spectra obtained by the EPIC cameras over a relatively broad energy band (0.15--12.0 keV) emphasize the various energy distributions observed from the different types of X-ray sources. A first analysis of twenty selected sources, presented here, yielded the discovery of a previously undetected supersoft source, revealed two ROSAT sources with indication for spatial extent as supernova remnants and rendered the detection of a possible pulsation in the X-ray flux of a candidate HMXB in the LMC. This increases the number of known HMXBs in the observed field to four. Among the fourteen brightest sources three HMXBs, one foreground star and up to ten AGN were found. Particular results on individual objects are discussed in Sect.\\,3 while in Sect\\,4 the analyzed sample of X-ray sources is considered in the view of source population studies in the Magellanic Clouds and nearby galaxies in general."
    },
    "0212/astro-ph0212543_arXiv.txt": {
        "abstract": "We present optical and infrared photometry of the unusual Type Ia supernova 2000cx. With the data of Li et al. (2001) and Jha (2002), this comprises the largest dataset ever assembled for a Type Ia SN, more than 600 points in $UBVRIJHK$.  We confirm the finding of Li et al. regarding the unusually blue $B-V$ colors as SN 2000cx entered the nebular phase.  Its $I$-band secondary hump was extremely weak given its $B$-band decline rate.  The $V$ {\\em minus} near infrared colors likewise do not match loci based on other slowly declining Type Ia SNe, though $V-K$ is the least ``abnormal''.  In several ways SN 2000cx resembles other slow decliners, given its $B$-band {\\em decline} rate ($\\Delta$m$_{15}$($B$) = 0.93), the appearance of Fe III lines and weakness of Si II in its pre-maximum spectrum, the $V-K$ colors and post-maximum $V-H$ colors. If the distance modulus derived from Surface Brightness Fluctuations of the host galaxy is correct, we find that the rate of light increase prior to maximum, the characteristics of the bolometric light curve, and the implied absolute magnitude at maximum are all consistent with a sub-luminous object with $\\Delta$m$_{15}$($B$) $\\approx$ 1.6-1.7 having a higher than normal kinetic energy. ",
        "introduction": "\\parindent = 9 mm  Astronomers try to understand the universe by looking for patterns in observed phenomena.  Often, the patterns themselves are reason to believe in underlying, understandable physical mechanisms, while at other times the exceptions to the rules provide the motivation to expand our conceptions of the physical makeup of cosmic objects. In this paper we present optical and infrared photometry of the very unusual supernova 2000cx. Previous optical data have been presented by Li et al. (2001), who describe SN 2000cx as ``unique''.  SN 2000cx was discovered by Yu, Modjaz, \\& Li (2000) from images taken on 17.5 and 18.4 July 2000 UT as part of the Lick Observatory Supernova Search, using the 0.76-m Katzman Automatic Imaging Telescope (KAIT).  This object was located at $\\alpha$ = 1:24:46.15, $\\delta$ = +9\\arcdeg ~30\\arcmin ~30\\farcs9 (equinox 2000.0), which is 23\\farcs0 west and 109\\farcs3 south of the nucleus of the S0 galaxy NGC 524. A spectrum taken on 23 July UT with the Nickel 1-m reflector at Lick Observatory (Chornock et al. 2000) revealed that the object was a peculiar Type Ia supernova, resembling SN 1991T a few days before maximum brightness, with prominent Fe III absorption lines near 430 and 490 nm but weak Si II at 612 nm. Optical photometry (Li et al. 2001)  revealed that SN 2000cx is different from all known Type Ia SNe and that the light curves cannot be fitted well using the techniques currently available.  The pre-maximum rise was relatively fast, similar to SN 1994D, but the post-maximum decline was relatively slow, similar to SN 1991T.  We present optical and infrared photometry of SN 2000cx, initiated at CTIO with the Yale-AURA-Lisbon-Ohio (YALO) 1-m telescope on 19 July 2000 (UT), some 8 days before the time of $B$-band maximum.  We include data taken with the 0.76-m Manastash Ridge Observatory (MRO) of the University of Washington (also begun on 19 July UT), the CTIO 0.9-m telescope, and the Apache Point Observatory 3.5-m telescope (APO). The calibration of the infrared light curves was primarily made possible with observations made with the Swope 1-m telescope at Las Campanas Observatory. ",
        "conclusions": "SN 2000cx, the brightest supernova discovered in the year 2000, occurred in the unobscured outer regions of an early-type galaxy and was well observed with multiple telescopes, allowing us to compile a datset of unprecedented size.  While many Type Ia SNe have light curves that follow patterns that are now well established, SN 2000cx did not conform to these patterns.  SN 2000cx was a reasonably fast riser in $B$ and $V$.  From the pre-maximum photometry Li et al. (2001) obtain a stretch factor that corresponds to $\\Delta$m$_{15}$($B$) = 1.64 $\\pm$ 0.02.  Based solely on its weak $I$-band secondary hump, we would have predicted $\\Delta$m$_{15}$($B$) $\\approx$ 1.7.  If the distance modulus based on Surface Brightness Fluctuations of the host galaxy is correct (Tonry et al. 2001), the corresponding absolute magnitudes in $BVIH$ are comparable to fast decliners, with $\\Delta$m$_{15}$($B$) in the range 1.4 to 1.7.  However, SN 2000cx was a slow {\\em decliner}, with $\\Delta$m$_{15}$($B$) = 0.93. Its pre-maximum spectrum showed strong Fe III and weak Si II, like other slow decliners such as SN 1991T.  Its $V-K$ color evolution, both pre- and post-maximum, was very similar to that of other slow decliners.  The bolometric behavior of SN~2000cx, when compared to the normal SNe 1999ee and 2001el, showed that this SN rose and fell from maximum light more rapidly, and that the magnitude difference between peak brightness and 90 days past peak was larger than normal. This behavior is consistent with the higher kinetic energies seen at maximum light and reported by Li et al. (2002), but it can also be explained by this event being sub-luminous.  The distance modulus of SN 2000cx is somewhat problematic. We note, however, that MLCS (Riess et al. 1996, 1998), Hubble Law's (with H$_0$ = 74 km s$^{-1}$ Mpc$^{-1}$), and the method that uses the $H$-band absolute magnitude at $t$ = 10 d (Krisciunas et al. 2003) give just about the same value.  The SBF method gives a distance modulus roughly 0.6 mag smaller.  Given: 1) the accuracy of the photometry of SN~2000cx at maximum ($\\pm$ 0.03 mag); 2) the host extinction of SN~2000cx is minimal (or zero); 3) the uncertainty of the Galactic extinction correction is also small; and 4) the light curves cannot be fit by templates based on other objects, confidently placing this SN in a Hubble diagram depends significantly on a direct measure of the host galaxy's distance (such as with the SBF method) or on an absolute magnitude derived from an explosion model that can match the many unusual observed facts."
    },
    "0212/astro-ph0212069_arXiv.txt": {
        "abstract": "{ Based on a  detailed study of the temperature structure of the intracluster medium in the halo of M~87, abundance profiles of 7 elements, O, Mg, Si, S, Ar, Ca, and Fe are derived. In addition, abundance ratios are derived from the ratios of line strengths, whose temperature dependences are small within the temperature range of the ICM of M~87. The abundances of Si, S, Ar, Ca and Fe show  strong decreasing gradients outside 2$'$ and become nearly constant  within the radius at $\\sim1.5$ solar. The Fe/Si ratio is determined to be 0.9 solar with no radial gradient. In contrast, the O abundance is less than a half of the Si abundance at the center and has a  flatter gradient. The Mg abundance is $\\sim$1 solar within 2$'$, which is close to stellar abundance within the same radius. The O/Si/Fe pattern of M~87 is located at the simple extension  of that of Galactic stars. The observed Mg/O ratio is about 1.25 solar, which is also the same ratio as for Galactic stars. The O/Si/Fe ratio indicates that the SN Ia contribution to Si and Fe becomes important towards the center and  SN Ia products have similar abundances of Si and Fe at least around M~87, which may reflect dimmer SN Ia observed in old stellar systems. The S abundance is similar to the Si abundance at the center, but has a steeper gradient. This result suggests that the S/Si ratio of SN II products is much smaller than the solar ratio. ",
        "introduction": "The intracluster medium (ICM) contains a large amount of metals, which are mainly synthesized in early-type galaxies (e.g. Arnaud et al. 1992; Renzini et al. 1993). Thus,  abundances of the metals are  tracers of chemical evolution of galaxies and clusters of galaxies.  Based on the  Si/Fe ratio observed with ASCA,  a discussion on  contributions from SN Ia and SN II to the metals has commenced. In a previous  nucleosynthesis model of SN Ia, the Fe abundance is much larger than the Si abundance in the ejecta of SN Ia (Nomoto et al. 1984). Observations of metal poor Galactic stars indicate that average products of SN II have a factor of 2--3 larger abundance of $\\alpha$-elements than Fe (e.g. Edvardsson et al. 1993;  Nissen et al. 1994; Thielemann et al. 1996), although this ratio may depend on the initial mass function (IMF) of stars. From elemental abundance ratios in the ICM of four clusters of galaxies observed with ASCA, Mushotzky et al. (1996) suggested that the metals in the ICM were mainly produced by SN II. Since stars in early-type galaxies in galaxy clusters  are very old (e.g. Stanford et al. 1995; Kodama et al. 1998),  this means that most of the metals in the ICM are produced through  star formations at high $z$. Fukazawa et al. (1998) systematically studied 40 nearby clusters and found that the Si/Fe ratio is lower among the low-temperature clusters, which indicates that SNe Ia products are also important among these clusters. From the observed radial dependence of the abundances, Finoguenov et al. (2000) found that the SNe II ejecta have been widely distributed in the ICM.  In the ICM, around a cD galaxy, the contribution of metals from the galaxy becomes important. Fukazawa et al. (2000) found a central increment of Fe and Si abundances around cD galaxies, which is  due to SNe Ia from the cD galaxies. The  supply of metals by stellar mass loss must be also considered (e.g. Matsushita et al. 2000)  In addition to the Si and Fe abundances, the XMM-Newton observatory enables us to obtain $\\alpha$-element abundances such as for O and Mg, which are not synthesized by SN Ia. B\\\"ohringer et al. (2001) and Finoguenov et al. (2002) analyzed annular spectra of M~87 observed by  XMM-Newton, and found a flatter abundance gradient of O  compared to  steep gradients of   Si, S, Ar,  Ca, and Fe. The stronger abundance increase of Fe compared to that of O indicates an enhanced SN Ia contribution in the central regions (in disagreement with our findings, Gastaldello and Molendi (2002) claim a similar gradient for O and Fe, which would not support this conclusion). The  observed similar abundance gradients of Fe and Si and the large Si/O ratio at the center imply a significant contribution of Si by SN Ia that is a larger Si/Fe abundance ratio by SN Ia than obtained by the classical model of Nomoto et al. (1984). The larger Si/Fe ratio and the implication that the Si/Fe ratio supplied by  SN Ia may change with radius in the M~87 halo may indicate a diversity of SN Ia explosions also reflected in a diversity of the light curves as discussed by Finoguenov et al. (2002). A similar abundance pattern of O, Si and Fe is observed around center of A 496 (Tamura et al. 2001).  A prerequisite of the abundance determination is a precise knowledge of the temperature structure of the ICM. Based on the XMM-Newton observation of M~87, Matsushita et al. (2002; hereafter Paper I) found that  the intracluster medium  has a single phase structure locally, except for the regions associated with  radio jets and lobes (B\\\"ohringer et al. 1995; Belsole et al. 2001), where there is an additional $\\sim$ 1 keV temperature component. The signature of gas cooling below 0.8 keV to zero temperature is not observed as expected for a cooling flow (e.g Fabian et al. 1984). The fact that the thermal structure of the intracluster medium is fairly simple and the plasma is almost locally isothermal facilitates the abundance determination enormously and helps to make the spectral modeling almost unique.  In this paper, based on the detailed study of the temperature structure, abundances of various elements are discussed. In section 2, we summarize the observation and data preperation. Sections 3 and 4 deal with the MEKAL (Mewe et al. 1995, 1996; Kaastra 1992; Liedahl et al. 1994) and the APEC model (Smith et al. 2001) fits to the deprojected spectra, and in section 5, abundance ratios are determined directly from the line ratios, considering their temperature dependences. In section 6, we evaluate an effect of  resonant line scattering. Section 7 gives a discussion of the obtained results, and Section 8 summarizes the paper. We adopt for the solar abundances the values given by Feldman (1992), where the solar Ar and Fe abundances relative to H are 4.47$\\times10^{-6}$ and 3.24$\\times10^{-5}$  in number, respectively. These values are different from the ``photospheric'' values of Anders \\& Grevesse (1989), where the solar  Ar and Fe abundances are 3.63$\\times10^{-6}$ and 4.68$\\times10^{-5}$, respectively. The solar abundances of the other elements are consistent with each other. Unless otherwise specified,  we use 90\\% confidence error regions. ",
        "conclusions": "Based on the temperature structure studied in Paper I,  abundance profiles  of  7 elements of the ICM around M~87 are obtained using the deprojected spectra. We have discussed the use of one- and two-temperature models in the fit and the application of the MEKAL and APEC codes. Two temperature models are only important in the region, R$<2$ arcmin. The previous problems with an earlier version of the APEC code are now resolved in the present version and this code provides now better fitting solutions in general. In addition, using line ratio profiles of the projected and deprojected data, the abundance ratio profiles are obtained with smaller uncertainties, considering the temperature dependence of the line ratio, The main results obtained are as follows,  \\begin{itemize} \\item Abundance profiles of 7 elements are obtained with high accuracy (e.g. $\\delta$Fe$\\sim\\delta$Si$\\sim$5\\%,  $\\delta$O$\\sim$10\\%). \\item       The Si and Fe abundance profiles have steep gradients at $>2'$ with a Fe/Si ratio of $0.9\\pm0.1$  solar, and become constant within this radius at 1.5--1.7 solar. \\item       The S/Si ratio is about 1 solar at $r<$ 4$'$ and decreases to $\\sim$ 0.7 solar at $r>10'$. \\item       The Ar/Si and Ca/Si are about 0.8 solar and 1.5 solar, respectively, although errors are relatively large. \\item  The O and Mg abundances are smaller than the abundances of Fe and the  intermediate elements.  The O/Si ratio is less than  half solar at the center and increases with radius. The Mg/O ratio is 1.25 solar and consistent with no radial gradients at least at $r>0.5'$. \\item The observed abundance pattern among O, Mg, Si, Ca and Fe of the ICM is located at an extension of that of Galactic stars. Therefore, the abundance pattern of the ICM is not peculiar and it should strongly constrain the products of SN Ia and SN II of early-type and late-type galaxies. \\item  The Si/O/Fe diagnostics shows large Si contributions by SN Ia, which may reflect the observational finding that SN Ia in old stellar systems are fainter. \\item  The observed Mg abundance of the ICM is consistent with the stellar metallicity profile from the Mg$_2$ index at the same radius. This result is consistent with stellar mass loss as the source of the gas in the very central region of M~87. \\item  The central abundance can be explained with the observed SN Ia rate and SN Ia model yields. \\item  The different radial profiles between Si and S suggest that the S/Si ratio synthesized from SN II is much smaller than the solar ratio. Then, the S/Si ratio may be a  better indicator of the relative contribution from SN Ia and SN II, when O and Mg abundances cannot  be observed. \\end{itemize}   We are actually very fortunate to have M~87 and the Virgo cluster as the closest galaxy cluster in our neighborhood, since it offers an ICM in just this low temperature range where the spectra are richest in spectral lines for this type of abundance diagnostics. Therefore it would be worthwhile to use this unique case for even deeper observations and more detailed investigations by X-ray spectroscopy in the future."
    },
    "0212/astro-ph0212296_arXiv.txt": {
        "abstract": "This paper describes a series of 3D simulations of shallow inefficient convection in the outer layers of the Sun. The computational domain is  a closed box containing the convection-radiation transition layer, located at the top of the solar convection zone. The most salient features of the simulations  are that: i)The position of the lower boundary can have a major effect on the characteristics of solar surface convection (thermal structure, kinetic energy and  turbulent pressure). ii)The width of the box has only a minor effect on the thermal structure, but a more significant effect on the dynamics (rms velocities). iii)Between the surface and a depth of 1 Mm, even though the density and pressure increase by an order of magnitude, the  vertical correlation length of vertical velocity is always close to 600 km. iv)  In this region the vertical velocity cannot be scaled by the pressure or the density scale height. This casts doubt on the applicability of the mixing length theory, not only in the superadiabatic layer, but also in the adjacent underlying layers. v) The final statistically steady state  is not strictly dependent on the initial atmospheric stratification. ",
        "introduction": "It is now just possible to perform physically realistic 3D simulations of the surface layers of the Sun which take into account the complex interaction between radiative and convective energy transports (Stein \\& Nordlund 1998, Kim \\& Chan 1998  hereafter denoted SN and KC, respectively, Stein \\& Nordlund 2000). There have also been a number of  2D simulations of the surface layers (Steffen et al. 1990  and Gadun et al. 2000). To model solar convection realistically, requires a realistic equation of state, realistic opacities and proper  treatment of radiative transfer in the shallow layers. SN and KC are the two most freqently cited 3D models of this type of stratified convection. As their  approaches differ considerably, both in numerical methods and in input model physics, it is important to find out, what particular aspect of the respective simulations, caused their results to be different.  The aim of this paper is to describe solar surface convection which not only has the KC realistic physics, but also has a realistic (as is presently possible) geometry. To achieve  this we had to increase both the depth and width of the original  KC model, until the side walls and the horizontal boundaries had only a minimal effect on the flow. The simulations themselves model a region less than a few thousand kilometers in depth, as measured inwards from the visible solar surface. In the deep regions of the solar convection zone, the turbulent velocity is subsonic and the superadiabaticity is close to zero. However, within a few hundred   kilometers  of the solar surface, the convective  flux starts to decrease. As the total flux is constant, the radiative flux must increase to offset the drop in the convective flux. This is achieved by a rise in the local temperature gradient $\\na$. This region of inefficient convection is called the superadiabatic layer (SAL).  In the SAL, the superadiabaticity $\\na -\\na_{\\rm ad}$ is positive and of order unity (Demarque, Guenther \\& Kim 1997, 1999). As the  buoyancy force is large in the SAL, the region is characterised by highly turbulent velocities and large relative thermodynamic fluctuations. The turbulent velocity also gives rise to a significant turbulent pressure (approximately 15 \\% of the gas pressure). This  moves out the convection surface, modifying the SAL and the stratification. In one-dimensional (1D) models of the solar convection zone based on the mixing length theory (MLT)(B\\\"ohm-Vitense 1958), the  velocity is set to zero above the convection boundary. However, three-dimensional (3D) numerical simulations described in Cattaneo et al. (1990), have shown that just above the top of the convection layer the turbulent velocities are still high.   There are several motivations for such  simulations  among stellar physicists.  One is to understand the effects of turbulence on the structure of the outer solar layers as revealed by the observed frequencies of solar {\\it p}-modes.  Another is to explain the excitation mechanism of the {\\it p}-modes. Still another is to derive more realistic surface boundaries for stars with convection zones from first physical principles, free of the arbitrary assumptions of the mixing length theory. And finally, such simulations may be of help in investigations of the solar dynamo.  Over the last few years, the science of helioseismology has  provided some very precise  measurements of the {\\it p}-mode oscillation frequencies   (to within one part in a thousand) in the surface layers of the Sun \\cite{harvey1996}. The discrepancy between the observed {\\it p}-mode frequencies and those calculated from solar models, is known to be primarily due to the approximations made in modelling the surface layers, where turbulent and radiative losses are significant (Balmforth 1992; Guenther1994). In the 3D simulation described in  KC, the turbulent pressure pushed the convective boundary radially outwards from its original position (that was computed using the MLT). This situation  was mimicked in the 1D solar model by tweaking  the opacity in the outer layers  \\cite{demarque1999}. This resulted in improved {\\it p}-mode frequencies for low and intermediate degrees. However, Demarque et al. also showed that the mixing length prescription of Canuto \\& Mazzitelli (1991), which had a completely different SAL structure than   KC, could produce a similar improvement of the {\\it p}-mode frequencies for the same $l$-values. Full details of the different approaches, the contrasting  SAL structures and the resulting {\\it p}-modes, are described in Demarque et al. (1997, 1999). Later, Rosenthal et al (1999)  used another approach to compute the {\\it p}-mode frequencies. These authors patched the mean stratification (horizontal average) of a 3D hydrodynamical simulation, onto a 1D MLT envelope model. In order to get a smooth 1D model, they adjusted the mixing length and the amplitude of the turbulent pressure to match the 3D simulation. The computed frequencies were closer to the observed frequencies than if a standard solar model is used, thus showing the importance of including turbulence in modeling the outer layers of the Sun. More recently, Li et al. (2002) found similar results to Rosenthal et al. by inserting the averaged turbulent pressure and turbulent kinetic energy directly into the 1D stellar model. Their method is applied to two of the simulations described in this paper (see section 4.3). As the turbulent kinetic energy and turbulent pressure were very small at the base of the 3D model, they were  set to zero in regions of the 1D stellar model that lay below the 3D model domain. This required the usual adjustment of the mixing length parameter and the helium abundance to calibrate the perturbed stellar model. No other adjustable parameters were employed. The improvement in the eigenfrequencies was found to be primarily  due to the inclusion of turbulent kinetic energy flux in the 1D stellar model.  After extensive testing, we found that our simulations are in good agreement with other numerical studies  of the surface layers (e.g. Rosenthal et al. 1999, Asplund et al. 2000). In addition, by incorporating the computed 3D turbulence into a  1D stellar model, we were able  to produce solar surface eigenfrequencies ({\\it p}-modes) that were very close to the observed frequencies. As our eventual goal is to simulate the SAL in stars other than the Sun, it is vital to be sure we are modelling the Sun as correctly as possible. ",
        "conclusions": "As we intend  to use our code to model the SAL in other stars, this paper is an important  benchmark for future  studies. From these simulations of the Sun, we found that: \\begin{enumerate} \\item The impenetrable lower boundary needs to be far enough away from the SAL so that by the time the fluid reaches the boundary, the velocities have become both weak and uncorrelated from velocities in the SAL. If the lower boundary is too close to the SAL, the kinetic energy will be overestimated. \\item The horizontal cross section needs to  be large enough so that the side-walls do not restrict the movement of the granules. If the box width is too small then the kinetic energy will be underestimated. \\item There is a region close to the surface in which the vertical correlation length remains constant, even though the density and pressure vary by an order of magnitude. \\item In that region the mixing length theory, which assumes the correlation length to be a constant multiple of the pressure scale height, will not work. \\item The final equilibrium state is not strictly dependent on the initial model atmosphere (see appendix, section \\ref{initial}). \\end{enumerate}  While these results  have  been found in a simulation of the Sun, it is reasonable to assume that similar criteria would apply to  the convection-radiation transition layers in other stars.  The effect of the boundaries should certainly be considered when simulating other convection-radiation  layers. The correlation length (half-width) seems to be  a very robust feature of the solar granules and could easily be measured in other computations. For example, in a recent proceedings (Robinson et al 2002) we describe the application of this  model to the Sun at the sub-giant and the start of the red giant branch. Preliminary results indicate that, near the surface of each convection zone, the  ratio of the half-width to the stellar radius  depends directly on the surface gravity.   \\vspace{3mm} This research is supported in part by NASA grant  NAG5-8406 (FJR, PD, SS). YCK is supported by a Department of Astronomy, Yonsei University grant (2001-1-0134). DBG's research is supported in part by a grant from NSERC of Canada. We also would like to acknowledge computer support from the  Centre for High Performance Computing at Saint Mary's University in Canada. Finally, we  thank H.-G. Ludwig  for helpful comments."
    },
    "0212/hep-ph0212240_arXiv.txt": {
        "abstract": "In the light of KamLAND data just released, we reanalyze and update the constraints on neutrino masses and mixing parameters, the most general ones that can be derived in three-flavor mixing scheme of neutrinos with use of the bounds imposed by neutrinoless double beta decay search and reactor experiments. We point out that with KamLAND data and assuming Majorana neutrinos one can derive, for the first time, an upper bound on neutrino contribution to the cosmological $\\Omega$ parameter, $\\Omega_{\\nu} \\leq 0.070 h^{-2}$, by using the current upper bound on mass parameters obtained by Heidelberg-Moscow group, $\\langle m \\rangle_{\\beta \\beta} < 0.35$ eV (90 \\% CL). If the bound is tighten to $\\langle m \\rangle_{\\beta \\beta} < 0.1$ eV by future experiments, it would lead to a far severer bound $\\Omega_{\\nu} \\lsim 0.01 h^{-2}$. ",
        "introduction": "The solar neutrino problem is finally solved by the KamLAND experiment \\cite {KamLAND}, which observed a clear deficit of about 40 \\% in $\\bar{\\nu}_e$ flux from reactors located in all over Japan. It pin-points the Large-Mixing-Angle (LMA) region as the unique solution of the solar neutrino problem based on the MSW matter enhanced neutrino flavor transformation \\cite{MSW}, excluding the remaining LOW and the vacuum (VAC) solutions. While the LMA solution was preferred over others on statistical ground in combined analyses of solar neutrino data \\cite{solar} it is by no means true that it was selected out as the unique solution. It is because the values of $\\chi^2$ of the LOW and the VAC solutions, $\\sim$ unity per degrees of freedom\\cite{solaranalysis}, were acceptable and hence it was not possible with only solar neutrino data to uniquely idetify the LMA solution.  It implies, first of all, that the three-flavor mixing scheme of neutrinos had further support by the experiments, and the structure of lepton flavor mixing is established together with the existing evidences which come from the atmospheric \\cite{SKatm} and the long-baseline accelerator experiments \\cite{K2K}. Given our above understanding of the KamLAND result it also means that the experiment gave the first direct proof of (almost) pure vacuum neutrino oscillation, the fascinating phenomenon first predicted by Maki, Nakagawa, and Sakata \\cite{MNS}, and first applied to the solar neutrino problem by Pontecorvo \\cite{pontecorvo}.   The resolution of the solar neutrino problem implies that we now know within certain uncertainties the values of $\\Delta m^2_{12}$ and $\\sin^2{2\\theta_{12}}$. It is an urgent task to uncover the all possible implications of the far more constrained values of these mixing parameters. We address in this paper the impact of the KamLAND result on masses and possible mass patterns of neutrinos.   We re-analyze the constraints on neutrino masses and mixing parameters imposed by neutrinoless double beta decay experiments in the light of the first results of KamLAND experiment under the assumption of Majorana neutrinos. See ref.~\\cite{vogel} for the present status of the double beta decay experiments. The constraints we will derive \\cite {mina-hiro02} are the maximal ones that can be derived in a model-independent way (and assuming ignorance of the CP violating phases) in a general framework of three-flavor mixing scheme of neutrinos. It utilizes the bounds imposed by neutrinoless double beta decay search and the CHOOZ reactor experiment \\cite{CHOOZ}.   A natural question would be \"what is the impact of the KamLAND data on the double beta decay bound on neutrino mass and mixing parameters?\". The answer is that it allows us to derive, for the first time, an upper bound on the neutrino contribution to the cosmological $\\Omega$ parameter, the energy density normalized by the critical density in the present universe. It is not because the parameter region for the LMA solution becomes narrower (in fact, the region for $\\theta_{12}$ is virtually identical with that before KamLAND), but because the LOW and the VAC solutions which allow maximal or almost maximal mixing are eliminated by the KamLAND data at 99.95~\\% CL \\cite{KamLAND}. ",
        "conclusions": "In the light of the KamLAND data we reanalyzed the constraint on neutrino masses and mixing imposed by neutrinoless double beta decay search and the CHOOZ reactor experiment. We have shown that by now the absolute mass scale for neutrinos is bounded from above by the ``conspiracy'' of KamLAND and double beta decay experiments. This conclusion and our whole analysis in this paper are valid under the assumption that neutrinos are Majorana particles.  In closer detail, allowed regions on a plane spanned by the lowest neutrino mass versus the solar mixing angle $\\theta_{12}$ were obtained by using $\\mbb^\\ma$ ($\\mbb^\\mi$), the experimental upper (lower) bound obtained and to be obtained in the ongoing and the future double beta decay experiments. For given $\\theta_{12}$, these become constraints on neutrino mass, the lightest mass $m_l$ (or the heaviest one though we did not give them explicitly). Roughly speaking, $\\mbb^\\mi \\leq m_l \\leq 3 [6]\\times\\mbb^\\ma$ for the LMA best fit parameters and for the most conservative case in square parenthesis, respectively.  It was made clear that the condition $|\\cos{2\\theta_{12}}| > t_\\CH^2 \\simeq 0.03$ was necessary for the upper bound on $m_l$ to exist. On the other hand, the condition $\\mbb^\\mi \\geq 0.0056\\,\\eV (0.048\\,\\eV)$ needs to be satisfied for the normal (inverted) mass hierarchy, for the lower bound on $m_l$ to exist. For example, $0.05\\,\\eV \\le \\mbb \\le 0.84\\,\\eV$ gives $0.05\\,\\eV (0.005\\,\\eV) \\leq m_l \\leq 2.5 [5.0]\\,\\eV$ for the normal (inverted) hierarchy.   \\vskip0.5cm  \\noindent {\\bf Note added}:  While this paper was in the revision process, we became aware of the report by the Wilkinson Microwave Anisotropy Probe (WMAP) \\cite{WMAP} in which a severe constraint on omega parameter $\\Omega_{\\nu} h^2 < 0.0076$ (95 \\% CL) is placed. It means that double beta decay searches must have the sensitivity in the region $\\mbb \\lsim 0.1\\,\\eV$ to be competitive, as one can see from Fig.~2."
    },
    "0212/astro-ph0212539_arXiv.txt": {
        "abstract": "Significant gravitational wave emission is expected from gamma-ray bursts arising from compact stellar mergers, and possibly also from bursts associated with fast-rotating massive stellar core collapses. These models have in common a high angular rotation rate, and observations provide evidence for jet collimation of the photon emission, with properties depending on the polar angle. Here we consider the gravitational wave emission and its polarization as a function of angle which is expected from such sources. We discuss possible correlations between the burst photon luminosity, or the delay between gravitational wave bursts and X-ray flashes, and the polarization degree of the gravitational waves. ",
        "introduction": "\\label{sec:intro}  The observations of gamma-ray burst (GRB) afterglows at energies ranging from X-rays to radio have lead to an increased understanding of the possible geometry of the inferred relativistic outflow or ejecta (see van Paradijs, Kouveliotou \\& Wijers 2000 for a review; also Rhoads 1999; Sari, Piran \\& Halpern 1999). At least for the class of long bursts (durations $\\gtrsim$ 10 s) there  is now significant evidence that the ejecta is collimated in a jet.  The evidence for jets is based on a change of the GRB light curve time-dependence power law index, and there appears to be a correlation in the sense that bursts with the largest gamma-ray fluences  have the narrowest implied opening angles. Frail et al. (2001), Piran et al. (2001) and Panaitescu \\& Kumar (2002) reported that the total gamma-ray energy release, after correcting for the collimation and distance determined by the afterglow observations, are narrowly clustered around $5 \\times 10^{50}$erg, and suggested that the broad distribution in fluence and luminosity for GRBs is largely the result of a wide variation of the opening angles. This result can be obtained assuming that the jet emission inside some angle $\\theta_j$ is approximately uniform and then drops off abruptly beyond that. However, Rossi, Lazzati \\& Rees (2002) and Zhang \\& M\\'{e}sz\\'{a}ros (2002) showed that the observations can also be interpreted in terms of an alternative model where the jet, rather than having a uniform profile out to some definite cone angle, has a universal angle-dependent beam pattern with a luminosity per unit solid angle which is maximal along the axis, and drops off gradually away from the axis. In this model, it is the difference of viewing angles $\\theta$ which causes a wide variation in the apparent luminosity of GRBs, $L \\propto \\theta^{-2}$.  Gravitational waves (GW) are also expected from some types of GRB (e.g. Kobayashi \\& M\\'{e}sz\\'{a}ros 2002). The GW emissivity and its polarization are angle-dependent, and may in principle be measurable, depending on the signal to noise ratio and on the detector alignments. In this Letter we discuss the prospects of exploiting measurements of the gravitational polarization degree to obtain information on the nature and orientation of the GRB system, providing additional constraints on the luminosity of the bursts, and we mention possible implications for the interpretation of dark GRB and X-ray flashes. ",
        "conclusions": ""
    },
    "0212/astro-ph0212013_arXiv.txt": {
        "abstract": "The observables in a strong gravitational lens are usually just the image positions and sometimes the flux ratios. We develop a new and simple algorithm which allows a set of models to be fitted exactly to the observations. Taking our cue from the strong body of evidence that early-type galaxies are close to isothermal, we assume that the lens is scale-free with a flat rotation curve. External shear can be easily included. Our algorithm allows full flexibility as regards the angular structure of the lensing potential. Importantly, all the free parameters enter linearly into the model and so {\\it the lens and flux ratio equations can always be solved by straightforward matrix inversion}. The models are only restricted by the fact that the surface mass density must be positive.  We use this new algorithm to examine some of the claims made for anomalous flux ratios. It has been argued that such anomalies betray the presence of substantial amounts of substructure in the lensing galaxy.  We demonstrate by explicit construction that some of the lens systems for which substructure has been claimed can be well-fit by smooth lens models. This is especially the case when the systematic errors in the flux ratios (caused by microlensing or differential extinction) are taken into account. However, there is certainly one system (B\\,1422+231) for which the existing smooth models are definitely inadequate and for which substructure may be implicated.  Within a few tens of kpc of the lensing galaxy centre, dynamical friction and tidal disruption are known to be very efficient at dissolving any substructure. Very little substructure is projected within the Einstein radius.  The numbers of strong lenses for which substructure is currently being claimed may be so large that this contradicts rather than supports cold dark matter theories. ",
        "introduction": "The fitting of models to observational data has always been a major concern in strong gravitational lensing (e.g., Schechter 2000). In most cases, the lens is only composed of 2 (a ``doublet'') or 4 (a ``quadruplet'') point-like images. For the doublets, the observable constraints are the four relative coordinates of the images with respect to the lensing galaxy, together with the one flux ratio of the first image to the second. For the quadruplets, this becomes eight relative coordinates and three flux ratios. Only for a handful of lens systems are time delays available.  It therefore happens very frequently that lens models have more degrees of freedom than the number of constraining observations.  In fact, the situation is often even worse than we have just described. Sometimes the centre of the lensing galaxy itself cannot be reliably identified. Sometimes a lens system is complicated by the possible effects of nearby bright galaxies. Very often, the flux ratios are untrustworthy, either because of differential extinction in the optical bands (e.g., Falco et al. 1999) or because of microlensing in the optical and radio (e.g., Irwin et al. 1989; Koopmans \\& de Bruyn 2000) or because of scintillation, scatter-broadening and free-free absorption in the radio (e.g., Patnaik et al. 1992, Jones et al. 1996, Winn, Rusin \\& Kochanek 2003). It is therefore already clear that considerable caution must be exercised in interpreting the results of fits to gravitational lens systems.  The approach of modellers has been largely two-pronged. First, general non-parametric methods have been developed (e.g., Saha \\& Wiliams 1997). These have proved powerful in exploring the range of degeneracies in the lensing mass that can give rise to a particular image configuration. For example, in the case of PG\\,1115+080, Saha \\& Williams present three different models obeying the lensing constraints, but for which the lensing mass distribution resembles a face-on spiral, an edge-on disk or a flattened elliptical galaxy respectively. Second, many simple parametric models have been thoroughly explored, such as those based on elliptically stratified potentials (e.g., Witt 1996; Witt \\& Mao 1997) or densities (e.g., Kassiola \\& Kovner 1993; Mu{\\~n}oz, Kochanek \\& Keeton 2001). The advantage of this is that the models may already incorporate some of the known properties of nearby galaxies.  However, this approach may also be dangerous because simple ansatze like elliptical potentials or surface mass densities can introduce unexpected properties. These properties may be so severe that, for example, the flux ratios may not be well fitted and wrong conclusions may be drawn. This can be seen most clearly in the case of elliptical potentials, in which there are strong constraints on the flux ratios (e.g., Witt \\& Mao 2000; Hunter \\& Evans 2001).  All this attains added significance in the light of recent claims of evidence for substructure from ``anomalous flux ratios'' in strong lensing. Mao \\& Schneider (1998) were the first to point out that substructure may be needed to explain the flux ratios in some cases. The instance that they selected, the quadruplet B\\,1422+231, has three highly magnified bright images and one much fainter image.  Three highly magnified images occur generically near a cusp, and a Taylor expansion gives a universal relationship that the sum of the fluxes of the two outer images should equal the flux of the middle image. This is strongly violated in B\\,1422+231 leaving substructure on top of the smooth model as the believable culprit. This result is supported by the detailed models of B\\,1422+231 by Bradac et al. (2002) which find that the discrepancy requires substructure on the mass scale $\\sim 10^6 \\msun$.  Recently, Dalal \\& Kochanek (2002) looked at seven predominantly radio quadruplets and argued that the flux ratios were anomalous by comparison with those expected for simple isothermal lenses with external shear.  They claimed that the anomalous flux ratios implied the existence of $\\sim 2$ per cent of the mass of the lensing galaxy in substructure.  Metcalf \\& Zhao (2002) looked at five optical quadruplets and similarly claimed evidence for anomalous flux ratios on comparison with those expected for simple elliptical power-law potentials with shear.  The problem with this procedure is, of course, that anomalous flux ratios may not be the result of substructure at all, but may simply reflect deficiencies in the modelling.  This motivates us to introduce in sections 2 and 3 a new approach for fitting which can incorporate the flux ratios at outset and which can permit an arbitrary azimuthal dependence for the surface density. Importantly, all the free parameters enter linearly into the models and so {\\it the lens and flux ratio equations can always be solved by straightforward matrix inversion}. The models are only restricted by the fact that the surface mass density must be positive. This algorithm is used in section 4 to assess the evidence for anomalous flux ratios. Readers who are mainly interested in this application may skip section 3 entirely on the first reading. This section is rather mathematical and explicitly derives the linear equations for the parameters to be fitted.  \\begin{figure*} \\epsfysize=7.5cm \\centerline{\\epsfbox{fig1.ps}} \\caption{The two plots show truncated Fourier series approximations to the equidensity contours of an exactly elliptical E3 (left panel) and an E7 (right) galaxy. For moderate flattenings such as E3, excellent results are obtained with only the first three non-vanishing terms in the expansion. Even for the most highly flattened configurations such as E7, the approximation becomes very accurate when the first seven non-vanishing terms are included.  } \\label{fig:sample} \\end{figure*} ",
        "conclusions": "The main aim of this paper is to present a new, simple approach to fitting a model to the data on a gravitational lens. Given the image positions and the flux ratios (if available), a model that exactly fits that data can be constructed by a simple matrix inversion (preferably by singular value decomposition, available in standard numerical libraries). All the fitting parameters enter linearly into the equation. No complicated non-linear $\\chi^2$-fitting needs to be done.  The lensing galaxy is always assumed to be scale-free with a flat rotation curve. There is already a substantial body of evidence from fitting of lenses that early-type galaxies have nearly isothermal profiles (Kochanek 1995; Mu{\\~n}oz et al. 2001; Cohn et al. 2001; Rusin et al. 2002). In particular, the absence of central images strongly constrains the lensing potential of early-type galaxies to be isothermal or steeper and the size of any core region to be small (e.g., Rusin \\& Ma 2001; Evans \\& Hunter 2002).  There is both stellar dynamical (e.g., Gerhard et al. 2001) and X-ray (Fabbiano 1997) evidence that early-type galaxies have flat rotation curves out to at least 4 effective radii. Hence, our assumption of a flat rotation curve seems exactly what the data require.  Unusually, our algorithm allows full flexibility as regards the angular structure of the lensing potential. Earlier fitting procedures have allowed the lensing potential to have different radial structure (for example, different power-law profiles), but usually only simple forms of angular structure (for example, constant external shear). In our models, the radial structure is always fixed as isothermal, but the shape of the density contours is completely arbitrary.  The lens model has a flexible number of free parameters, and of course includes the isothermal elliptic density and potential models which have a distinguished history in gravitational lens modelling.  The fit allows direct deduction of the mass contour of the lensing galaxy.  Although we have presented examples based on quadruplet lens systems, the algorithm can be adapted for a two image or a higher image model.  A higher image model might be needed, for example, in Q\\,0957+561, where 10 sub-image positions can be detected in the radio band. Another good application might be MG\\,0414+0534, where some substructure of the images are detected (cf. Trotter, Winn \\& Hewitt 2000 and their fitting approach).  Both lenses seem heavily distorted by external shear. A shear term can be added to our algorithm. If the shear is known, then the lens and flux ratio equations remain linear.  The flux ratios may enter directly into the fit depending on whether they are available or trustworthy.  We stress that flux ratios often do not provide good model constraints.  The reason for this is easy to understand. The flux ratios depend on the second derivatives of the lensing potential, whereas time delays and image positions are proportional to the potential and its first derivative respectively. Therefore, flux ratios are particularly sensitive to graininess in the gravitational potential. This often manifests itself as microlensing. The flux ratios are also affected by the uncertain differential extinction corrections that must be applied to each image.  Radio and mid-infrared fluxes are generally more reliable than optical.  However, effects such as scintillation, scatter-broadening and free-free absorption may affect the radio fluxes in some of the lenses. If flux ratios are incorporated into a fit, it is vital that there is a realistic treatment of the errors.  Our fitting algorithm permits the construction of a whole class of degenerate models all of which can fit the image position and flux ratios exactly.  In addition, the models are also degenerate in the time delay since the delay depends only on the distance of the image positions to the centre of the lensing galaxy (cf. Witt, Mao \\& Keeton 2000).  Such degeneracy is easy to understand pictorially. Given a Fermat time delay surface, we need only keep the values, the derivatives and the second derivatives of the surface at the image positions fixed. We then have enormous freedom to move the surface (subject only to the constraints that no additional images are introduced and that no negative mass density is produced).  We have used our new fitting algorithm to examine critically some of the claims made recently for anomalous flux ratios. The procedure used by both Dalal \\& Kochanek (2002) and Metcalf \\& Zhao (2002) was to show that a family of simple models did not reproduce the flux ratios. However, it does not follow from this that substructure is necessary; it may just be that the simple model did not have enough flexibility to provide a good match. For example, of the five lenses studied by Metcalf \\& Zhao (2002), they themselves concluded that one (MG\\,0414+0534) could be satisfactorily explained by a smooth model. In this paper, we have demonstrated that the data on two others (PG\\,1115+080, Q\\,2237+030) are consistent with smooth galaxy-like models, especially when a realistic treatment of the errors in the flux ratios is incorporated.  We note that the fraction of mass in substructure expected in galaxy haloes is very small ($\\lta 5$ per cent) and it occurs overwhelmingly in the outlying portions (e.g., Moore et al. 1999; Ghigna et al. 2000). Substructure evolves as it is subjected to tides, impulsive collisions and dynamical friction. Tidal disruption becomes important as soon as the mean density of the galaxy is equal to the density of the substructure. Similarly, the dynamical friction timescale scales with the square of the galactocentric radius (see e.g., Binney \\& Tremaine 1987). So, both tides and dynamical friction are efficient at erasing substructure in the inner parts. This effect is referred to as ``anti-biasing'' by Ghigna et al. (2000). So, on dynamical grounds, little substructure is projected within the Einstein radius.  Taking the example of PG\\,1115+080, the projected Einstein radius is just $\\sim 3.6$ kpc. The fraction of mass in substructure projected within a cylinder of radius $\\sim 3.6$ kpc is clearly much lower than the global fraction of $5$ per cent, which pertains to the entire dark halo of total extent $\\sim 200$ kpc. It is crucial that lensing calculations do not assume an everywhere uniform fraction of substructure in the halo, as this does not take into account the ``anti-biased'' spatial distribution of substructure and therefore necessarily over-emphasises the importance of the effects of substructure on flux ratios. Notice that an interesting consequence of the spatial distribution is that anomalous flux ratios are more likely to occur for lenses with larger angular separation, as these have larger Einstein radii.  The real question at issue is the following. Suppose a simple lens model (such as a simple isothermal ellipsoid plus shear) does not adequately fit the data on positions and flux ratios. What can be legitimately deduced? The difficulty is that there are many modifications of the simple model that remove the discrepancy. One of these is substructure, as pointed out by Dalal \\& Kochanek and Metcalf \\& Zhao.  As shown in this paper, another is higher order multipoles in the lensing mass, such as diskiness, boxiness, lopsidedness and barredness (see also M\\\"oller, Hewett \\& Blain 2003, who make a similar point).  It is crucial to develop techniques to distinguish between higher order quadrupoles on the one hand and substructure on the other.  From our modelling, we tend to agree that the substructure candidate B\\,1422+231 originally pointed out by Mao \\& Schneider (1998) is strong.  Using a novel application of the cusp relation, Keeton, Gaudi \\& Petters (2002) have argued that B\\,2045+265 and RX J\\,0911+0551 may be two further good candidates.  It is an outstanding problem to predict -- for different cosmologies -- how many quadruplets may have anomalous flux ratios. Accurate calculations will become possible only when distributions of mass and position of substructure become available from high resolution simulations. It will be interesting to see whether the numbers of lenses for which substructure is currently being claimed are compatible with cold dark matter theories."
    },
    "0212/astro-ph0212155_arXiv.txt": {
        "abstract": "{Compressible turbulence, especially the magnetized version of it, traditionally has a bad reputation with researchers. However, recent progress in theoretical understanding of incompressible MHD as well as that in computational capabilities enabled researchers to obtain scaling relations for compressible MHD turbulence. We discuss scalings of Alfven, fast, and slow modes in both magnetically dominated (low $\\beta$) and gas pressure dominated (high $\\beta$) plasmas. We also show that the new regime of MHD turbulence below viscous cutoff reported earlier for incompressible flows persists for compressible turbulence. Our recent results show that this leads to density fluctuations. New understanding of MHD turbulence is likely to influence many key astrophysical problems. } ",
        "introduction": " ",
        "conclusions": "In this paper we have discussed the new outlook onto compressible MHD turbulence. Contrary to common beliefs the compressible MHD turbulence does not present a complete mess, but demonstrates nice scaling relations for its modes. A peculiar feature is that those relations should be studied locally, i.e. in the frame related to the local magnetic field. However, such a system of reference is natural for many phenomena, e.g. for cosmic ray propagation. Recent application of the scalings obtained for compressible turbulence have shown that fundamental revisions are necessary for the field of high energy astrophysics. For instance, Yan \\& Lazarian (2002) demonstrated that fast modes dominate cosmic ray scattering even in spite of the fact that they are subjected to collisional and collisionless damping. This entails consequences for models of cosmic ray propagation, acceleration, elemental abundances etc.  Advances in understanding of MHD turbulence have very broad astrophysical implications. The fields affected span from accretion disks and stars to the ISM and the intergalactic medium in clusters. Turbulence is known to hold the key to many astrophysical processes. It was considered too messy by many researchers who consciously or subconsciously tried to avoid dealing with it. Others, more brave types, used Kolmogorov scalings for compressible strongly magnetized gas, although they did understand that those relations could not be true. \\adjustfinalcols Recent research in the field provides the scaling relations and insights that will contribute to many areas of research."
    },
    "0212/astro-ph0212363_arXiv.txt": {
        "abstract": "We determine the linear amplitude of mass fluctuations in the universe, $\\sigma_8$, from the abundance of massive clusters at redshifts $z=0.5-0.8$. The evolution of massive clusters depends exponentially on the amplitude of mass fluctuations and thus provides a powerful measure of this important cosmological parameter. The relatively high abundance of massive clusters observed at $z>0.5$, and the relatively slow evolution of their abundance with time, suggest a high amplitude of mass fluctuations: $\\sigma_8 =0.9\\pm 10$\\% for $\\Omega_m =0.4$, increasing slightly to $\\sigma_8 =0.95$ for $\\Omega_m =0.25$ and $\\sigma_8 =1.0$ for $\\Omega_m =0.1$ (flat CDM models). We use the cluster abundance observed at $z=0.5-0.8$ to derive a normalization relation from the high-redshift clusters, which is only weakly dependent on $\\Omega_m$: $\\sigma_8 \\Omega_m^{0.14}=0.78\\pm 0.08$. When combined with recent constraints from the present-day cluster mass function, $\\sigma_8 \\Omega_m^{0.6}=0.33\\pm 0.03$, we find $\\sigma_8 =0.98\\pm 0.1$ and $\\Omega_m =0.17\\pm 0.05$. Low $\\sigma_8$ values ($\\la 0.7$) are unlikely; they produce an order of magnitude fewer massive clusters than observed. ",
        "introduction": "\\label{sec:intro}  The amplitude of mass fluctuations is a fundamental cosmological parameter that describes the normalization of the linear spectrum of mass fluctuations in the early universe -- the spectrum that seeded galaxy formation. The abundance of massive clusters depends exponentially on this parameter (assuming Gaussian initial fluctuations), because a high amplitude of mass fluctuations forms structure rapidly at early times, while a lower amplitude forms structure more slowly.  The most massive systems ($\\sim 10^{15}h^{-1}M_\\odot$), which take the longest time to form and grow, did not exist at early times if the initial amplitude of mass fluctuations is low, but rather formed only recently.  The amplitude parameter, denoted $\\sigma_8$ when referring to the {\\it rms} linear density fluctuation in spheres of radius $8h^{-1}$Mpc at $z=0$, is not easily determined since the mass distribution cannot be directly observed. As a result, this parameter is not yet accurately known. Recent observations suggest an amplitude that ranges in value from  $\\sigma_8\\sim 0.7$ to a `high' value of $\\sigma_8\\sim 0.9-1$. While the difference in the reported values is only around 50\\%, the impact on structure formation and evolution is much larger, since the latter depends exponentially on $\\sigma_8^2$. The low amplitude values are suggested by current observations of the CMB spectrum of fluctuations \\citep{netterfield, sievers02, bondea02, ruhlea02}. However, this $\\sigma_8$ determination is degenerate with the unknown optical depth at reionization: if the optical depth were underestimated, then $\\sigma_8$ would be as well\\footnote{See note at end of paper}. Recent observations of the present-day cluster abundance as well as cosmic shear lensing measurements have also suggested that $\\sigma_8\\sim 0.7$ \\citep[e.g.][]{jarvis03, hamana02, sel01}. However, these measures provide a degenerate relation between the amplitude $\\sigma_8$ and the mass-density parameter $\\Omega_m$: $\\sigma_8 \\Omega_m^{0.6}\\approx 0.33$ \\citep{ike02, SDSSmf03, jarvis03, sel01}. The amplitude $\\sigma_8\\sim 0.7$ is implied only if $\\Omega_m \\sim 0.3$. If $\\Omega_m \\sim 0.2$, as is suggested by some observations \\citep[e.g.][]{CYE97, bah98, bah00, wklc01, ike02, rei02}, then the amplitude is $\\sigma_8\\sim 0.9-1$.  Early results from the Sloan Digital Sky Survey (SDSS) cluster data \\citep{SDSSmf03} use the shape of the observed cluster mass function to break the degeneracy between the parameters and find $\\sigma_8 =0.9^{+0.3}_{-0.2}$ and $\\Omega_m =0.19^{+0.08}_{-0.07}$. Similar results have recently been obtained from the temperature function of a large sample of X-ray clusters \\citep{ike02, rei02}. Most of the recent cluster normalization observations, as well as cosmic shear lensing measurements suggest $\\sigma_8\\simeq 0.9-1$ if $\\Omega_m \\simeq 0.2$ \\citep[][and the references above]{jarvis03, hamana02}. Combining current CMB measurements with the SDSS cluster mass function yields intermediate values of $\\sigma_8 =0.76\\pm 0.09$ and $\\Omega_m =0.26^{+0.06}_{-0.07}$ \\citep{MBBS03}.   The evolution of cluster abundance with time, especially for the most massive clusters, breaks the degeneracy between $\\sigma_8$ and $\\Omega_m$ \\citep[e.g.][]{PDJ89, ECF96, OB97, BFC97, CMYE97, bah98, DV99, Henry00}. This evolution depends strongly on $\\sigma_8$, and only weakly on $\\Omega_m$ or other parameters. The expected abundance of massive clusters at $z\\sim 0.5-1$ differs between Gaussian models with $\\sigma_8=0.6$ and $\\sigma_8=1$ by orders-of-magnitude, nearly independently of other parameters \\citep[e.g.]{FBC97}. Therefore, this method provides a uniquely powerful tool in estimating the amplitude $\\sigma_8$.  In this paper we use the abundance of the most massive clusters observed at $z\\sim 0.5-0.8$  to place a strong limit on $\\sigma_8$. ",
        "conclusions": "\\label{sec:disc}  We use the observed abundance of high-mass clusters of galaxies at $z=0.5-0.8$ to determine the linear amplitude of mass fluctuations, $\\sigma_8$. The cluster abundance depends exponentially on this amplitude, and only weakly on other parameters; it therefore provides a powerful method for measuring this important parameter. We show that the relatively high abundance of massive clusters observed at $z \\gtrsim 0.5$, as well as their relatively slow evolution with time, requires a high amplitude of mass fluctuations, $\\sigma_8 \\sim 0.9-1$. This conclusion is nearly independent of the exact value of $\\Omega_m$ (in the typical range of $\\Omega_m \\sim 0.1-0.4$).  We use the observed abundance at $z\\gtrsim 0.5$ to determine a normalization relation from high redshift clusters. The relation depends only weakly on $\\Omega_m$: $\\sigma_8\\Omega_m^{0.14} = 0.78\\pm 0.08$; alternatively, a linear relation of the form $\\sigma_8 =1.03-0.3\\Omega_m$ ($\\pm$10\\%) provides a similarly good fit to the data. These fits illustrate that $\\sigma_8 \\gtrsim 0.8$ for any $\\Omega_m \\leq 0.4$. For the typical observationally suggested value of $\\Omega_m \\simeq 0.2-0.3$, the amplitude is $\\sigma_8 =0.95\\pm 0.1$. We emphasize that this high $\\sigma_8$ value indicated by the cluster abundance at high redshift is nearly independent of the exact value of $\\Omega_m$.  We combine the high redshift constraint above with the independent normalization relation obtained from low redshift cluster abundance--- a relation that is steeper in $\\Omega_m$ ($\\sim \\Omega_m^{0.6}$; equation [\\ref{eqn:sdssfit}]). The combination breaks the degeneracy between the two parameters. We find $\\sigma_8 =0.98\\pm 0.1$ and $\\Omega_m =0.17\\pm 0.05$ (Figure \\ref{fig:hizont}; equation [\\ref{eqn:bestpar}]).  The high value of $\\sigma_8$ required to explain the high abundance of the most massive clusters at $z \\gtrsim 0.5$ is consistent with the present day cluster mass function if the mass density parameter is low, $\\Omega_m \\sim 0.2$. If $\\Omega_m$=0.3, then the high redshift clusters still require a high amplitude ($\\sigma_8 =0.92 \\pm 0.1$), since this constraint is nearly independent of $\\Omega_m$; the low redshift cluster abundance is consistent with this value at the 2-sigma level.  These results improve upon \\citet{bah98} by using improvements over the standard Press-Schechter formula (which does not accurately reproduce results from cosmological simulations at high redshift), allowing for changes in $h$ and $n$, and by using a more recent cluster normalization relation at low redshift.  The current results yield slightly lower values for the cosmological parameters than the previous work (the latter suggested $\\sigma_8 =1.2\\pm 0.22$ and $\\Omega_m =0.2^{+0.13}_{-0.07}$ (68\\%) when using the most massive clusters) but are consistent with the new values within the error-bars \\citep{bah98,FBC97}.   The excess CMB fluctuations detected on small scales by the CBI \\citep{mason02} and the BIMA \\citep{dawsea02} experiments implies (if correctly interpreted as being due to the S-Z effect from distant clusters) that $\\sigma_8 =1.04 \\pm 0.12$ \\citep[95\\%;][]{komsel02}. This is in excellent agreement with our current conclusions. We note, however, that this high amplitude is inconsistent with the lower value of $\\sigma_8\\approx 0.7$ suggested by current CMB data on large scales (which is degenerate with the unknown optical depth). Future CMB observations should clarify this current inconsistency. If massive clusters exist with relatively high abundance at high redshifts, as suggested by the data used here \\citep[as well as by deep X-ray surveys, e.g.][]{RBN02}, then these clusters should indeed produce the excess S-Z fluctuations observed by the CMB data.  The relatively high abundance of massive clusters observed at $z\\gtrsim 0.5$ provides one of the strongest arguments for a high amplitude of mass fluctuations, $\\sigma_8\\simeq 1$.  {\\bf Note added March 7, 2003:} The recent CMB anisotropy spectrum released by the WMAP team in February 2003 nicely confirms the results presented here. The constraint from the CMB alone is $\\sigma_8 =0.9\\pm 0.1$ \\citep{SpergWMAP}, in full agreement with the current high redshift cluster constraint."
    },
    "0212/astro-ph0212424_arXiv.txt": {
        "abstract": "{ We describe an automatic procedure for determining abundances from high resolution spectra. Such procedures are becoming increasingly important as large amounts of data are delivered from 8m telescopes and their high-multiplexing fiber facilities, such as FLAMES on ESO-VLT. The present procedure is specifically targeted for the analysis of spectra of giants in the Sgr dSph; however, the procedure may be, in principle, tailored to analyse stars of any type. Emphasis is placed on the algorithms and on the stability of the method; the external accuracy rests, ultimately, on the reliability of the theoretical models (model-atmospheres, synthetic spectra) used to interpret the data. Comparison of the  results of the procedure with the results of a traditional analysis for 12 Sgr giants shows that abundances accurate at the level of 0.2 dex, comparable with that of traditional analysis of the same spectra, may be derived in a fast and efficient way. Such automatic procedures are not meant to replace the traditional abundance analysis, but as an aid to extract rapidly a good deal of the information contained in the spectra. ",
        "introduction": "In recent years considerable attention has been devoted to full or partial automation of the process of abundance determination from high resolution spectra \\citep{Katz98,KatzT,EN02,EN02b}. One of the main thrusts behind these attempts is the increasingly large amount of data delivered by modern instrumentation. It is well known  that astronomical archives are full of good quality high resolution spectra which have not been analysed, or only partially analysed, due to lack of manpower. In spite of the complexity of data reduction, this is not the bottleneck, thanks to the efficient software and instrument-dedicated pipelines that are available. The real bottleneck is the data analysis which, for abundances, is still done more or less in the same way as twenty years ago. The situation is going to become even more critical as high-multiplex instruments, such as FLAMES on ESO-VLT \\citep{Pasquini}, become fully operative.  In the last years we have devoted special interest to the determination of abundances in giants of the Sgr dSph \\citep{B99,bonivlt,B00}. Such an astrophysical problem is ideal for the capabilities of FLAMES. It has also the advantage that two of the key parameters in the interpretation of stellar spectra, effective temperature and surface gravity, may be conveniently constrained. If we pick stars of roughly the same apparent magnitude, if they belong to Sgr, they will have the same luminosity and therefore surface gravity. At a given luminosity the RGB of Sgr spans a limited range in effective temperature, in fact less than 1000 K, and in the case of our sample less than 250 K. We had already attempted to develop a procedure to analyse the low resolution spectra which we obtained from EMMI-NTT \\citep{B99,bonivlt}, however the procedure, based on spectral indices, provided unsatisfactory results, essentially because of the combination of low resolution and low S/N ratios. Furthermore it was apparent that at low resolution we had no handle on the microturbulent velocity, and since the abundances relied only on strong lines (the only ones available at low resolution, whichever the S/N ratio), the result was strongly dependent on the unknown microturbulent velocity. Having in mind the future use of FLAMES and having available a few UVES spectra, obtained in slit-mode, we decided to concentrate on a procedure capable of analysing spectra from UVES and Giraffe. We have developed a procedure which is very efficient and stable, at least on the real UVES spectra. In spite of the fact that it is highly targeted (it will deal only with stars of a given luminosity and a limited range of effective temperatures), it has been written in such a way that it may be modified to deal with other types of stars. In this paper we describe the procedure placing our emphasis on the algorithms and on the stability of the method. The procedure rests on synthetic spectra computed from 1D LTE model atmospheres. In this paper we do not question the reliability of the input synthetic spectra, since it is straightforward to replace the existing grid with a better one, when available.   \\begin{table} \\caption{Algorithm meta-code} \\label{algorithm} \\begin{center} \\begin{tabular}{lll} \\hline \\\\ \\\\ BEGIN &\\\\ & $\\xi$;\\\\ &&  pseudo-normalize;\\\\ &&  find [Fe/H] for each FeI feature;\\\\ &&  find slope A(FeI) vs.  EW;\\\\ &&  if slope$<$ threshold, go to $\\alpha$;\\\\ &&  find a $\\xi$ to make slope smaller;\\\\ & goto  $\\xi$;\\\\ & $\\alpha$;\\\\ &&  find  [Mg/H] for each Mg I  feature;\\\\ &&  find  [Ca/H] for each Ca I feature;\\\\ && [$\\alpha$/Fe] = mean of [Mg/Fe] and [Ca/Fe];\\\\ && if  $\\Delta$[$\\alpha$/Fe] $<$ threshold1 goto END;\\\\ & goto $\\xi$;\\\\ END   \\\\    \\\\ \\hline \\\\ \\end{tabular} \\end{center} \\end{table}    \\begin{figure} \\centering \\resizebox{\\hsize}{!}{\\includegraphics[clip=true]{histomet.eps}} \\caption{Histogram of the derived [Fe/H] for Monte Carlo simulations with different S/N ratios and a resolving power of 7 kms$^{-1}$} \\label{histo} \\end{figure}  \\begin{figure} \\centering \\resizebox{\\hsize}{!}{\\includegraphics[clip=true]{fig_mc.eps}} \\caption{ Monte Carlo simulations corresponding to a resolving power of 7 kms$^{-1}$: [Fe/H], [$\\alpha$/Fe] and microturbulence velocity as a function of S/N for input spectrum at $\\xi =1.$Km/s, [$\\alpha $/Fe]=0.4 and [Fe/H]=--0.5 on the left panel and [Fe/H]=--1.5 on the right panel. For both the [Fe/H] vs S/N plots the smaller error bars correspond to Monte Carlo simulation dispersion, the bigger one to the [Fe/H] dispersion over the all Fe\\,\\textsc{i} features.} \\label{snplot} \\end{figure}  \\begin{figure*} \\centering \\resizebox{16cm}{!}{\\includegraphics[clip=true]{code_vs_trad.eps}} \\caption{Comparison of [Fe/H], [$\\alpha$/Fe] and $\\xi$ from the traditional analysis and the automatic analysis for the 12 stars of Sgr observed with UVES. The left panels display  the results from the UVES spectra, i.e. at a resolving power of $\\sim 7\\rm kms^{-1}$; the right panels display  the results from the UVES spectra convolved with a gaussian profile, obtaining a resolving power of $\\sim 20 \\rm kms^{-1}$. The two points, in each of the eight panels,  shown as crossed squares, identify the two stars for which the log g adopted in the traditional analysis is not 2.5.} \\label{conf} \\end{figure*}  \\begin{figure} \\centering \\resizebox{8.5cm}{!}{\\includegraphics[clip=true]{pl_br7_br20.eps}} \\caption{Comparison of [Fe/H], [$\\alpha $/Fe] and $\\xi$ obtained for  the 12 Sgr giants observed with UVES directly from the UVES spectra (resolving power $\\sim \\rm 7 kms^{-1}$) and UVES spectra broadened to a resolving power of $\\sim \\rm  20kms^{-1} $.} \\label{br7_20} \\end{figure} ",
        "conclusions": ""
    },
    "0212/astro-ph0212048_arXiv.txt": {
        "abstract": "We present an analysis method that allows us to estimate the Galactic formation of radio pulsar populations based on their observed properties and our understanding of survey selection effects. More importantly, this method allows us to assign a statistical significance to such rate estimates and calculate the allowed ranges of values at various confidence levels. Here, we apply the method to the question of the double neutron star (NS--NS) coalescence rate using the current observed sample, and we find calculate the most likely value for the total Galactic coalescence rate to lie in the range $3-22$ Myr$^{-1}$, for different pulsar population models.  The corresponding range of expected detection rates of NS--NS inspiral are $(1-9) \\times 10^{-3}$ yr$^{-1}$ for the initial LIGO, and $6-50$ yr$^{-1}$ for the advanced LIGO. Based on this newly developed statistical method, we also calculate the probability distribution for the expected number of pulsars that could be observed by the Parkes Multibeam survey, when acceleration searches will alleviate the effects of Doppler smearing due to orbital motions.  We suggest that the Parkes survey will probably detect $1-2$ new binary pulsars like PSRs~B1913+16 and/or B1534+12. ",
        "introduction": "The detection of the double neutron star (NS--NS) prototype PSR~B1913+16 as a binary pulsar (Hulse \\& Taylor 1975) and its orbital decay due to emission of gravitational waves have inspired a number of quantitative estimates of the coalescence rate, $\\cal{R}$, of NS--NS binaries (Clark et al.\\ 1979; Narayan et al.\\ 1991; Phinney 1991; Curran \\& Lorimer 1995). Significant interest derives from their importance as gravitational-wave sources for the upcoming ground-based laser interferometers (such as LIGO).  We present a newly developed statistical analysis that allows the calculation of statistical confidence levels associated with rate estimates. The method can be applied to any radio pulsar population. Here, we consider PSR~B1913+16 (Hulse \\& Taylor 1975) and PSR~B1534+12 (Wolszczan 1991). For different assumed distributions of pulsar properties (luminosities, Galactic positions), we derive the probability distribution function of the total Galactic coalescence rate weighted by the two observed binary systems. The method involves the simulation of selection effects inherent in all relevant radio pulsar surveys and a Bayesian statistical analysis for the probability distribution of ${\\cal R}$. The small-number bias and the effect of the faint-end of the luminosity function, previously identified as the main sources of uncertainty in rate estimates (Kalogera et al.\\ 2001) are {\\it implicitly included} in this analysis. We extrapolate the Galactic rate to cover the detection volume of LIGO and estimate the most likely detection rates of NS--NS inspiral events for the initial and advanced LIGO. Details of this work are given in Kim et al.\\ (2003; hereafter KKL).  In the second part of this paper, we modify our statistical method in a way that allows us to calculate the probability distribution for the number of pulsars that could be detected by the Parkes Multibeam survey (hereafter PMB survey; Lyne et al.\\ 2000; Manchester et al.\\ 2001) PMB, when the effects of Doppler smearing due to orbital motions are corrected with acceleration searches. ",
        "conclusions": "We have recently developed a new method for estimating the total number of pulsars in our Galaxy and have applied it to the calculation of the coalescence rate of double neutron star systems in the Galactic field (for more details see KKL). Here, we extend this method to obtain a prediction for the average number of observed pulsars that the PMB survey could detect when acceleration searches are used to correct for the Doppler smearing due to orbital motions.  The modeling of pulsar survey selection effects is formulated in a ``forward'' way, by populating the Galaxy with model pulsar populations and calculating the likelihood of the real observed sample. This is in contrast to the ``inverse'' way of the calculation of scale factors used in previous studies.  We note that this method could be further extended to account for distributions of pulsar populations in pulse periods, widths, and orbital periods. It is important to note that both our rate estimates and the predictions for detections from the PMB survey do not apply to binary pulsars that are significantly different from with such properties that are significantly different PSRs~B1913+16 and B1534+12 in terms of pulse shapes and orbital properties.  Most importantly the method can be applied to any type of pulsar population with appropriate modifications of the modeling of survey selection effects. Currently we are working on assessing the contribution of double neutron stars formed in globular clusters as well as the formation rate of binary pulsars with white dwarf companions that are important for gravitational-wave detection by LISA, the space-based interferometer planned by NASA and ESA for the end of this decade."
    },
    "0212/astro-ph0212562_arXiv.txt": {
        "abstract": "We have determined the composite luminosity function (LF) for galaxies in 60 clusters from the 2dF Galaxy Redshift Survey. The LF spans the range $-22.5<M_{b_{\\rm J}}<-15$, and is well-fitted by a Schechter function with ${M_{b_{\\rm J}}}^{\\!\\!\\!\\!*}=-20.07\\pm0.07$ and $\\alpha=-1.28\\pm0.03$ ($H_0$=100\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_M$=0.3, $\\Omega_\\Lambda$=0.7). It differs significantly from the field LF of \\cite{mad02}, having a characteristic magnitude that is approximately 0.3~mag brighter and a faint-end slope that is approximately 0.1 steeper. There is no evidence for variations in the LF across a wide range of cluster properties: the LF is similar for clusters with high and low velocity dispersions, for rich and poor clusters, for clusters with different Bautz-Morgan types, and for clusters with and without substructure. The core regions of clusters differ from the outer parts, however, in having an excess of very bright galaxies. We also construct the LFs for early (quiescent), intermediate and late (star-forming) spectral types. We find that, as in the field, the LFs of earlier-type galaxies have brighter characteristic magnitudes and shallower faint-end slopes. However the LF of early-type galaxies in clusters is both brighter and steeper than its field counterpart, although the LF of late-type galaxies is very similar. The trend of faint-end slope with spectral type is therefore much less pronounced in clusters than in the field, explaining why variations in the mixture of types do not lead to significant differences in the cluster LFs. The differences between the field and cluster LFs for the various spectral types can be qualitatively explained by the suppression of star formation in the dense cluster environment, together with mergers to produce the brightest early-type galaxies. \\\\ \\\\ \\noindent {\\bf Key words:}\\quad galaxies: luminosity function, mass function -- galaxies: formation ",
        "introduction": "The galaxy luminosity function (hereafter LF), which describes the number of galaxies per unit volume as a function of luminosity, is a fundamental tool for testing theories of galaxy formation and interpreting observations of galaxies at high redshift for evidence of evolution. Furthermore, precise and accurate measurements of the LF in different environments have the potential to provide important clues as to the role of `environmental' processes (e.g., dynamical interactions in rich clusters) in determining the properties of the present-day galaxy population.  One of the main legacies of the numerous redshift surveys that have been undertaken in the last decade or more is the wealth of LF measurements for galaxies in the low density `field' environment. A compilation and comparison of these various measurements was recently published by \\citet{cro01}. This work showed that while the data were adequately represented by a Schechter function, there were serious discrepancies between the various measurements, with the LFs differing by as much as a factor of 2 at the $L^{*}$ point, and with a scatter of a factor of 10 at $0.01L^{*}$. These differences are due to a combination of surface brightness selection, colour, aperture effects and local density variations among others. These problems have now been overcome in the recent analysis of large redshift surveys such as the 2dF Galaxy Redshift Survey (2dFGRS: Folkes et al. 1999, Madgwick et al. 2001, Norberg et al. 2002) and the Sloan Digital Sky Survey (SDSS: Blanton et al. 2001).  Rich clusters of galaxies provide the other extreme in environment, representing the highest density regions inhabited by galaxies. It is generally easier to derive LFs in clusters, since they provide rich ensembles of galaxies all at the same distance, whose over-density with respect to the surrounding field is sufficiently high to efficiently identify members either photometrically, through the statistical removal of foreground and background galaxies (e.g., Driver et al. 1998a and references therein) or spectroscopically \\citep{sma97,dep98,ada00}  These techniques have been used to measure LFs for individual clusters, which in many cases have been combined to form a `composite' LF to improve statistics (particularly at the brightest luminosities) and average out systematic uncertainties \\citep{dre78, lug86, col89, lug89, gai97, lum97, val97, ram98, gma99, pao01, got02, yag02}. However, these studies have not been unanimous on the exact form of the LF, with some claiming there to be significant differences between the LFs from cluster to cluster and between cluster and field \\citep{dre78, lpc97, lum97, val97, gma99, dri98a, got02}, while others finding no differences and concluding that galaxies in all environments appear to be drawn from a single, `universal' LF \\citep{lug86, col89, lug89, gai97, ram98, tre98, pao01, yag02}. A summary of the Schechter function fits from some of these previous studies (all of which are based on the technique of background subtraction, unlike the present work which uses spectroscopic identifications of cluster members) is given in Table 1. We have transformed magnitudes to $H_0=100$ km s$^{-1}$ Mpc$^{-1}$ but we have not changed their cosmology.  \\begin{table*} \\begin{center} \\begin{minipage}{140mm} \\caption{Composite LFs for rich clusters in blue passbands.} \\begin{tabular}{ccccc} \\hline Reference & M$^*$ & $\\alpha$ & N$_{\\rm clusters}$ & Luminosity Range \\\\ \\hline Schechter (1976) & $-19.9 \\pm 0.5$ & $-1.24$ & 13 & $-22.5 < M_J < -18.5$ \\\\ Dressler (1978) & $-19.7 \\pm 0.5$ & $-1.25$ & 12 &$ -23.5 < M_F < -18.5$ \\\\ Colless (1989) & $-20.10 \\pm 0.07$ & $-1.25$ & 14 & $-22.5 < M_J < -17$ \\\\ Lumsden et al. (1997)\\footnote{Average of three methods} & $-20.16 \\pm 0.02$ & $-1.22 \\pm 0.04$ & 46 & $-21 < M_b < -18$ \\\\ Valotto et al. (1997) & $-20.0 \\pm 0.1$ & $-1.4 \\pm 0.1$ & 55 & $-21 < M_b < -17$ \\\\ Rauzy et al. (1998) & $-19.91 \\pm 0.21$ & $-1.50 \\pm 0.11$ & 28 & $-21 < M_b < -17$ \\\\ Garilli et al. (1999) & $-20.30 \\pm 0.10$ & $-0.94 \\pm 0.07$ & 65 & $-22.5 < M_g < -15.0$ \\\\ Paolillo et al. (2001) & $-20.22 \\pm 0.15$ & $-1.07 \\pm 0.08$ & 39 & $-24.5 < M_g < -16.5$ \\\\ Goto et al. (2002) & $-21.24 \\pm 0.11$ & $-1.00 \\pm 0.06$ & 204 & $-24 < M_{g'} < -17$ \\\\ \\hline \\end{tabular} \\end{minipage} \\end{center} \\end{table*}  With the 2dF Galaxy Redshift Survey (2dFGRS; Colless et al. 2001) now complete, there is an opportunity to revisit these issues and simultaneously address the detailed form of the LF in rich clusters and in the field in a consistent manner. High-quality measurements of the field LF based on 2dFGRS data have already been published by \\cite{mad02} and \\cite{nor02}. In this paper we present an analysis of the LFs of 60 rich clusters, taken from the sample of known clusters within the 2dFGRS survey region that were identified and characterised by \\cite{dep02}.  The enormity of 2dFGRS in terms of its size, depth and sky coverage has a number of distinct advantages in comparison to previous LF studies. Firstly, cluster membership is determined unambiguously from spectroscopic redshifts for nearly all galaxies, eliminating the introduction of systematic errors into the derived LFs through field subtraction \\citep{dri98b}. Secondly, the apparent magnitude limit of 2dFGRS ($b_{J}=19.45$) is sufficiently deep that, at the redshifts covered here ($z<0.11$), our study extends to $\\sim 5$ magnitudes fainter than $M^{*}$ -- at least as deep as previous studies. Thirdly, the sheer number of clusters (60) that we can study together with the almost 1-in-1 sampling of their galaxy populations, provides the level of statistical discrimination needed, particularly at the bright end of the LF \\citep{col89}, to detect differences that are of physical interest. Finally, our comparison of the cluster and field LFs is done entirely within the 2dFGRS and hence based on the same input catalogue, galaxy photometry and redshift observations. The field and cluster samples therefore share most of the selection effects and observational biases and the resulting LFs can be compared fairly: we also note that the field LF in \\cite{mad02} is derived from galaxies with $z < 0.15$ and is therefore similar to the volume-limited sample of cluster galaxies.  The plan of this paper is as follows: The next section describes our cluster sample and the procedure used for constructing composite LFs. We then present our derived LFs in Section 3, both for the entire sample of clusters and for subsets differentiated on the basis of velocity dispersion, Bautz-Morgan type, richness, and the presence of substructure. Section 4 compares our data with previous work and the field. Our results are discussed and summarised in Section 5. We adopt the `concordance' cosmology with $\\Omega_m=0.3$, $\\Omega_{\\Lambda} =0.7$ and $H_0=100$. This is about 0.07 magnitudes brighter than the Einstein-DeSitter model at the mean redshift ($z=0.07$) of our cluster sample. ",
        "conclusions": "We used the simulations of \\cite{col89} to analyze the ability of our data to detect differences in LF parameters: we find that our LFs are of sufficient quality to detect differences of 0.15 mag in $M^*$ and 0.07 in $\\alpha$ at the $1\\sigma$ level, which is an improvement by a factor of more than 2 over previous studies. Furthermore, unlike all previous work, we are not limited by uncertainties in background subtraction. In the following we compare our results to previous determinations of both the cluster and field LFs and study the universality of the LF with cluster subsamples. Type-specific LFs are also discussed. We finally consider the implications of our findings for models of galaxy formation.  \\subsection{Comparison with Previous Work}  A comparison with previous work is shown in Table 1. Among previous studies only \\cite{gma99}, \\cite{pao01} and \\cite{got02} reach luminosity limits comparable to ours. \\cite{gma99} derive $M^*_g=-20.30 \\pm 0.10$ and $\\alpha=-0.94 \\pm 0.07$. \\cite{pao01} derive $M^*_g=-20.22 \\pm 0.15$ and $\\alpha=-1.07 \\pm 0.08$ from their 39 clusters. Once we correct for the colour difference between $g$ and $b_{\\rm J}$ and for the different cosmologies (using the mean redshifts of the cluster sample) the characteristic luminosities of \\cite{gma99} and \\cite{pao01} agree with our value within their 2$\\sigma$ errors, whereas the two $\\alpha$s differ at about 2$\\sigma$; the two LFs agree with each other at about the 1.5$\\sigma$ level: the difference is only marginally significant.  The LF of \\cite{got02} from SDSS data is in considerable disagreement with ours: $M^*$is brighter by about 1.2 magnitudes and $\\alpha$ is flatter. This LF disagrees with \\cite{col89,lum97} and \\cite{val97}, all of which use APM data. It also disagrees with the $M^*$ and $\\alpha$ values of \\cite{gma99} and \\cite{pao01} in the $g$ band. One possibility is a systematic offset between APM and SDSS photometry, but this appears to be ruled out by a comparison of SDSS Early Release Data, the Millennium Galaxy Catalogue and 2dF photometry (Cross et al. 2002, in preparation).  Earlier work is generally more limited in luminosity coverage: values for the Schechter function fits for \\cite{col89, lum92, val97} and \\cite{ram98} are, once we apply the necessary cosmological and filter corrections, in reasonable agreement with our data. With the exception of \\cite{got02}, the agreement with previous work is generally satisfactory.  \\subsection{Cluster versus Field}  One of the motivations behind this work and previous studies is to test the universality of the LF and its dependence on the environment. The most dramatic comparison in this context is between rich clusters and the field, because of the factor of 100 or more difference in galaxy density between the two environments.  Having derived a statistically robust composite cluster LF from the 2dFGRS data, of foremost interest now is to make this comparison with its 2dFGRS counterpart for the field. We base this comparison on the LFs published by \\cite{mad02}; these cover the absolute magnitude range $-23.0 < M_{b_{\\rm J}} < -13.5$, and include an overall field LF, and LFs derived by dividing the galaxies into four different spectral classes (see Figure 1). The $M^{*}_{b_{\\rm J}}$ and $\\alpha$ parameter values for their Schechter function fits to these LFs are listed at the bottom of Table 3. We also compare these LFs directly using a two-sample $\\chi^2$ test, the results of which, for both $M_{b_{\\rm J}} < -16.0$ (the full magnitude range) and $M_{b_{\\rm J}} < -18$ (the brighter cluster members), are shown in Table 4.  \\begin{table} \\begin{center} \\caption{$\\chi^2$ comparisons of the field and cluster LFs.} \\begin{tabular}{ccc} \\hline Sample & P($\\chi^2$)         & P($\\chi^2$)         \\\\ & $M_{b_{\\rm J}}<-18$ & $M_{b_{\\rm J}}<-16$ \\\\ \\hline All      & $<10^{-3}$ & $<10^{-3}$ \\\\ Type 1   & $<10^{-3}$ & $<10^{-3}$ \\\\ Type 2   & 0.133      & $<10^{-3}$ \\\\ Type 3+4 & 0.556      & 0.320      \\\\ \\hline \\end{tabular} \\end{center} \\end{table}  The error ellipses for the Schechter function fits to all these LFs are compared in Fig.~5. We only show the 3$\\sigma$ contour for field galaxies as the errors are small.  \\begin{figure} \\includegraphics[width=84mm]{FIG6.ps} \\caption{ Error ellipses for the Schechter function fits to the field and cluster LFs. The total LFs are represented by solid lines (thin for clusters and thick for the field, as in Fig.~4). Dashed lines are for type 1 galaxies, dot-dashed lines for type 2 galaxies and dotted-dashed lines for types 3+4. We only show the 3$\\sigma$ error ellipse for the field data. } \\end{figure}  The $\\chi^2$ comparison shows that the overall cluster and field LFs differ at more than 3$\\sigma$, and this conclusion is also borne out by a comparison of the error contours presented in Fig.~5. In terms of the individual LF parameters, the overall cluster LF is 0.3 magnitudes brighter in $M^*$ and 0.1 steeper in $\\alpha$ than the overall field LF. We note here that although the `field' LF is actually the total LF for all environments, clusters do not contribute to the field sample to any great extent, as cluster members are only 3\\% of the total number of galaxies in the survey; even for bright galaxies this fraction is less than 6\\%.  Significant differences are also found when the field and cluster LFs for the different spectral classes are compared. For ease of comparison, Fig.~6 shows type-dependent LFs for both field and cluster galaxies; the the field LFs are normalized so that the overall field LF matches the overall cluster LF at $M_{b_{\\rm J}}=-18$.  \\begin{figure*} \\includegraphics[width=160mm]{FIG5.ps} \\caption{ Comparison of cluster and field LFs. The open symbols are the type-completeness corrected cluster counts shown in Figure 3. Thick lines represent the field LFs from Madgwick et~al.\\ (2002), normalized to agree with the cluster LF at $M_{b_{\\rm J}}=-19$. The two thin lines show the luminosity distribution for no completeness correction (lower line) and for the maximal correction discussed in the text (upper line). Error bars are excluded for clarity (see Fig.~3 for error bars). The error ellipses are shown in Fig.~5 for both the field and cluster samples. } \\end{figure*}  The $\\chi^2$ tests show that the LFs for type 1 cluster and field galaxies are different over their entire range. This is due to the fact that the type 1 LF in clusters is both brighter and much steeper than in the field. The LFs of type 2 galaxies also differ, but Fig.~6 both shows that the difference is primarily due to the steeper faint-end slope; the $\\chi^2$ fit shows that, to $M_{b_{\\rm J}} < -18$, the two LFs match adequately. Finally, the LF of star-forming galaxies (types 3 and 4) are essentially identical in clusters and the field. These differences in the shapes of the types-specific LFs, and the different contributions from each type (as shown in Fig.~1), account for the differences between the overall cluster and field LFs. The steeper faint-end slopes for type 1 and type 2 galaxies in clusters may lead, at magnitudes fainter than those covered by our sample, to an upturn of the effective LF slope, such as has been claimed to exist in some clusters \\citep{dep98b}.  One caveat concerns the assumption implicit in equation~9, that the galaxies without spectral types have the same type distribution as those for which types could be determined. In order to test the importance of this issue, we have computed LFs for each of the spectral types without any completeness correction. We find that $M^*$ is hardly changed whereas $\\alpha$ is flattened by about 0.1; nevertheless, the LFs of type 1 and 2 galaxies remain steeper than in the field whereas types 3+4 are consistent with the field LF. Incompleteness in spectral type comes from two sources: galaxies which were not observed and therefore have neither redshift nor spectral type and which are therefore an unbiased sample, and galaxies whose spectra have too low a signal-to-noise to yield a redshift or a spectral type. The untyped galaxies are roughly divided evenly between these two categories. It is only the latter one that may suffer from bias, in that galaxies of a specific type may be preferentially misidentified. We have therefore adopted a `maximal correction' where all of these galaxies are assigned to each type in turn. This effect is small for types 1 and 2 but may be significant for types 3 and 4. This is shown in Fig.~6. However, no matter what correction is applied, cluster galaxies of all types have LFs with faint-end slopes as steep or steeper than their counterparts in the field.  These differences between the type-dependent LFs in the cluster and field samples appear to be inconsistent with previous claims for a universal type-dependent LF \\citep{bst88,jeta97,and98}. However, it needs to be noted that these earlier studies are concerned with morphological types, which are only moderately well correlated with our spectral types (Fig.~1), so that a given spectral type may include a range of morphological types.  It is instructive at this stage to compare LF parameters for field and cluster galaxies, divided by spectral type. Fig.~7 shows that, progressing from late to early types, the cluster LFs have brighter $M^*$ and flatter $\\alpha$, as observed in the field LFs \\citep{mad02}. The trend in the value of $M^*$ with type is stronger in clusters than in the field, with a similar value of $M^*$ for type 3+4 galaxies but a significantly brighter $M^*$ for type 1 galaxies. However the trend of $\\alpha$ with type is weaker for cluster galaxies than for field galaxies---at faint magnitudes the type 3+4 LF is steeply rising ($\\alpha=-1.3$) in both clusters and the field, but the type 1 LF in the field is actually falling at faint magnitudes ($\\alpha=-0.52$) while in clusters it is merely flat ($\\alpha=-1.05$). This trend of $\\alpha$ with environment for type 1 galaxies is apparent even within the cluster sample: the type 1 galaxies in richer clusters have a LF with a steeper faint end than those in poorer clusters, which in turn have a steeper LF than type 1 galaxes in the field (see Table~3 and Fig.~7). This is consistent with the trend for star formation rate to be broadly anticorrelated with the faint-end slope, which has been observed by \\cite{mar02} for galaxies in groups identified within the 2dFGRS 100K release. Both \\cite{yag02} and \\cite{got02} claim similar variations in their samples.  \\begin{figure} \\includegraphics[width=84mm]{FIG7.ps} \\caption{ Comparison of cluster and field LF parameters. Open symbols are for field galaxies and filled symbols for cluster galaxies; also shown are the LF parameters for type 1 galaxies in richer clusters (triangle down) and poorer clusters (triangle up). } \\end{figure}  \\subsection{Cluster subsamples}  We have also compared subsamples of clusters chosen according to velocity dispersion, richness, Bautz-Morgan type and likelihood of containing substructure: similarly, we also considered samples of galaxies within and outside two King core radii from the cluster centre. These results are summarized in Table~3. We find that the values of $M^*$ and $\\alpha$ for each pair considered (e.g.\\ low and high velocity dispersion) do not differ at more than 1.5$\\sigma$ in all cases. The only possible exception is for the case of clusters containing substructure, where $\\alpha$ appears to be steeper. In all cases (again, excepting clusters with likely substructure) the derived LF parameters are consistent with those of the total LF within about 1.5$\\sigma$. This is in contrast with \\cite{lum97}, who provided weak evidence for differences in the LFs of clusters with high and low $\\sigma$, and with \\cite{val97}, who suggested differences between rich and poor clusters.  The difference in the LF of clusters with substructure is potentially interesting, as it is generally believed that substructure is an indicator of recent or ongoing cluster merging. However an examination of the confidence contours for the fitted Schechter function parameters, shown in Fig.~4, shows that when the correlated nature of the parameters is considered, the fits to the LFs of the clusters with and without substructure are in fact consistent at better than 2$\\sigma$.  The similarity of the LFs for the various subsamples is confirmed by two-sample $\\chi^2$ tests, which indicate that the probabilities that the pairs of contrasted subsamples have consistent LFs are: 99.3\\% for early and late Bautz-Morgan clusters, 35.6\\% for high and low $\\sigma$ clusters, 50.0\\% for rich and poor clusters, 27.1\\% for clusters with and without substructure and 0.9\\% for galaxies within and outside of 300 kpc from the cluster centre. Inspection of Fig.~4 shows that the LFs of galaxies in the inner 300 kpc and the outer regions of clusters differ in detail; in particular, the inner region LF is a poor fit to the Schechter function, with a deficit of $L^*$ galaxies and an excess of brighter objects. This is reminiscent of the galactic cannibalism scenario of \\cite{osha78} and \\cite{muri84}, where $L^*$ galaxies are preferentially destroyed to fuel the growth of giant ellipticals, D and cD galaxies.  An alternate approach to quantifying the differences between LFs is to calculate the dwarf to giant ratio (D/G). We define giants as galaxies with $-21.5 < M_{b_{\\rm J}} < -18$ and dwarfs as galaxies with $-18 < M_{b_{\\rm J}} < -16.0$, as in \\cite{dri98a}. These ratios are tabulated in Table~3 for all the LFs we consider. In general, our cluster samples have D/G ratios comparable with those of moderately rich systems in \\cite{dri98a}. The D/G ratios for the contrasting pairs of subsamples only differ at the $\\sim$1$\\sigma$ level, except for clusters with and without substructure (1.5$\\sigma$) and the inner and outer regions of clusters (2.2$\\sigma$)---in no case is the difference in D/G ratio highly significant.  The overall impression is therefore one of broad universality of the LF over a range of cluster properties. This is surprising if one considers the very considerable differences between the LFs of the various spectral types in the field \\citep{mad02}, since different mixtures of types would lead to different LFs. For instance, the morphology-density relation might lead one to expect that rich and massive clusters would be more elliptical-rich than poor, low-mass clusters; likewise, early B-M clusters should be dominated by ellipticals, while late B-M type clusters should be spiral-rich. Mixing the field LFs of the different morphological mixes in these differing proportions would lead to clusters with significantly different LFs. The reason this does not appear to be the case is that the LFs of the different types are more similar in clusters than they are in the field, so that changing the mixture of types within clusters has less effect than one might have expected.  \\subsection{Implications for Galaxy Formation}  The main conclusions of our analysis are as follows:  (i) We have determined the composite LF of galaxies in clusters from the 2dFGRS. The LF is well-fitted by a Schechter function with parameters $M^*_{b_{\\rm J}}=-20.07 \\pm 0.07$ and $\\alpha=-1.28 \\pm 0.03$. This is significantly different to the field LF of \\cite{mad02}, having a characteristic magnitude that is approximately 0.3~mag brighter and a faint-end power-law slope that is approximately 0.1 steeper.  (ii) There is no significant evidence for variations in the LF across a broad range of cluster properties; the LF appears similar for clusters with high and low velocity dispersions, for rich and poor clusters, for clusters with early and late Bautz-Morgan types, and for clusters with and without substructure. However the core regions of clusters differ from the outer parts in having an excess of very bright galaxies.  (iii) Breaking down the LF by spectral type, the same trends are apparent in clusters as in the field: the LFs of earlier-type galaxies, with lower star-formation rates, have brighter characteristic magnitudes and shallower faint-end slopes; the LFs of later-type galaxies, with higher star-formation rates, have fainter characteristic magnitudes and steeper faint-end slopes. The trend in faint-end slope, which is the dominant difference between the LFs of the various spectral types in the field, is much less pronounced in clusters. The smaller differences between the LFs of different spectral types in clusters explain why variations in cluster properties giving rise to significant variations in the mixture of types do not lead to significant differences in the cluster LFs.  (iv) A comparison between the field and cluster LFs of each spectral type reveals that while the LF of late-type, star-forming galaxies is very similar in clusters and the field, the LF of early-type galaxies with low star-formation rates is both brighter and steeper in clusters than in the field; intermediate types in clusters have an LF with a similar bright end but a steeper faint end than their field counterparts. As early and intermediate spectral types are predominant in clusters, the overall cluster LF is also brighter and steeper than the overall field LF.  The above results may be compared with two recent redshift-based studies of galaxy clusters. \\cite{dep98} derived the $K$-band LF of galaxies in the inner 25~arcmin of the Coma cluster from a 100\\%-complete sample of members, and found that the cluster's LF is indistinguishable from the general field LF. \\cite{cz02} use a redshift survey of 6 clusters (with a redshift completeness of 20--50\\%) to derive $R$-band LFs for both the cluster members and the field galaxies in the fore- and background of the clusters. They also find that the $R$-band field and cluster LFs are very similar. These results are for samples selected in red and infrared passbands, and are therefore less sensitive to star formation, to which our $b_{\\rm J}$-selected sample is more closely correlated. In contrast to these results, we find that there {\\em are} small but statistically significant differences between field and cluster galaxies, with the blue-selected cluster LF being both brighter and steeper than that in the field. The similarity between the field and cluster LFs in the red and near-infrared passbands suggests that field and cluster galaxies with the same total stellar mass have similar integrated star-formation histories, while the difference between the field and cluster LFs in the blue simply means that the star-formation rate is significantly affected by the cluster environment.  Up until now it has been possible to claim that the observed differences between cluster and field LFs were due in large part simply to the different proportions in which supposedly universal type-specific LFs were mixed in the two environments. This explanation is now precluded by the finding that the type-specific LFs differ significantly with environment (in fact, as Fig.~5 shows, the difference between clusters and the field is less for the overall LF than for some of the type-specific LFs).  Another way to understand the differences between the field and cluster LFs is to consider a simple `closed box' model. In this model one considers the type-specific field LFs to be the initial LFs within the volumes which today have collapsed to form the clusters. The relative normalizations of the LFs are based on the observed relative numbers of the different types in the field and in clusters (types 1:2:3+4 in the proportions 0.36:0.32:032 in the field and 0.54:0.24:0.22 in clusters; see Fig.~1) and the assumption that the total number of objects in the initial and final LFs is the same (i.e.\\ neglecting mergers within the closed-box volume). Evolution from the initial (field) LFs to the final (cluster) LFs occurs almost entirely through processes which suppress star-formation as the cluster collapses and becomes denser \\citep{lew02}, converting galaxies from later types to earlier types. In the simplest version of this model, this suppression of star-formation does not affect the galaxies' luminosities. Fig.~8 shows the initial (field) and final (cluster) LFs of each spectral type in this naive model, and compares these LFs to the observed cluster LFs.  \\begin{figure} \\includegraphics[width=84mm]{FIG8.ps} \\caption{ A comparison of the field LF (solid line) and closed-box model cluster LF (dotted line) to the observed cluster LF (dashed line) for each spectral type. Note that for type 3+4 the dotted line is identical to the dashed line.} \\end{figure}  This model has some successes: it is consistent with the fact that the type 3+4 cluster LF has the same shape but a lower normalization than the field LF, and it does cause the initial field LFs of the type 1 and 2 galaxies to evolve towards the forms they are observed to have in clusters. However this passive steepening of the LFs, achieved by shifting galaxies from the steeper LFs of later types to the shallower LFs of earlier types, fails to steepen the LFs sufficiently to reproduce the observed clusters.  A more successful model requires both passive steepening as galaxies shift from later to earlier types and also active steepening due to luminosity-dependent fading. In such a model, the suppression of star-formation affects about a third of type 3+4 galaxies, independent of luminosity. The affected galaxies' instantaneous star-formation rate decreases, lowering the strength of their $H\\alpha$ emission so that they become type 2 galaxies. This leaves the shape of the type 3+4 LF the same, but reduces the numbers of these strongly star-forming galaxies. When type 2 galaxies undergo further suppression of their star-formation rate, this has the effect of decreasing their $b_{\\rm J}$ luminosities and eventually converting them to type 1, where further fading may occur. This fading must be greater for brighter galaxies in order to actively steepen the LF slopes and reproduce the observed cluster LFs.  In addition, some merging (or cannibalism) is required to explain the small number of very bright type 1 galaxies in clusters which are not present in the field LFs, and may also help explain the deficit of bright type 2 and $L^*$ type 1 cluster galaxies. This is consistent with the conclusions of \\cite{cz02}, who find that the only difference between the $R$-band LFs of cluster and field galaxies is an excess of bright non-starforming galaxies. However the higher dwarf-to-giant ratio in the LFs of earlier-type galaxies in clusters indicates that star-formation suppression is more important than mergers in shaping the cluster LF.  It should be noted that the field LF is in fact the mean LF of the entire galaxy population, and is therefore dominated by galaxies belonging to groups. Thus our results therefore imply that suppression of star-formation is the dominant effect in evolving from the typical group environment to the rich cluster environment. Merger effects are expected to be more important in groups, however, and this will be investigated in a future paper based on the group catalogue derived directly from the 2dFGRS.  The closed-box model is over-simplified in a variety of ways, and can only serve as a qualitative guide to understanding the processes shaping the LF of cluster galaxies. Unfortunately, a quantitative interpretation of our results is hampered by the fact that few theoretical studies have considered the evolution of the LF and its dependence on environment, and most have been limited to the brighter cluster members (e.g.\\ Malumuth \\& Richstone 1984). Detailed semi-analytic models are clearly required to explore the relative importance of the various mechanisms that may be driving the evolution of galaxies in clusters. These models will need to be complemented by more stringent observational constraints to distinguish the effects of the different processes involved. It will be particularly important to obtain the near-UV and near-IR LFs for large spectroscopic samples of cluster galaxies. These will yield both the instantaneous star-formation rate and the total stellar mass of the galaxies, and reveal where galaxies are evolving due to mergers and where they are undergoing changes in their star-formation rate."
    },
    "0212/nucl-th0212056_arXiv.txt": {
        "abstract": "Using the Relativistic Mean Field Theory (RMF) it is shown that different proton fraction which is directly connected with the neutron excess and with the asymmetry of the system affects proto-neutron stars parameters and changes their composition. The obtained form of the equation of state allows to construct the mass-radius relations and shows that the increasing asymmetry creates more compact stars. The inclusion of $\\delta $ meson together with nonlinear vector meson interaction terms and magnetic field make this effect even stronger. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/astro-ph0212204_arXiv.txt": {
        "abstract": "We use time-resolved spectra of the cataclysmic variable RW Tri in the I and K-band to determine the orbital velocity of the secondary star using skew mapping and cross correlation techniques respectively. We find radial velocity amplitudes of $250\\pm47$ km/s in the I-band, and $221\\pm29$ km/s in the K-band. We also determine the rotational velocity of the secondary star using the K-band data and find a $V_{rot}\\sin i$ of $120\\pm20$km/s. A combination of these results coupled with an estimate of the effect of heating on the secondary star suggests a mass ratio $M_2/M_1$ in the range $0.6-1.1$; the mass ratio range with no correction for heating is $0.5 - 0.8$. These lead to most likely estimates of the primary and secondary star masses in the range $0.4-0.7M_{\\odot}$ and $0.3-0.4M_{\\odot}$ respectively. Further refinement of the stellar masses is hampered by uncertain knowledge of the white dwarf orbital velocity, and we discuss evidence that at least some estimates of the white dwarf velocity are contaminated by non-orbital components. ",
        "introduction": "RW Tri is an eclipsing nova-like cataclysmic variable (CV) system with an orbital period of 5 hours 34 mins.  Nova-like systems contain a non-magnetic white dwarf primary and a late type (K-M) secondary star that fills its Roche lobe, transferring mass to the white dwarf via an accretion disc.  Since first being observed as an eclipsing variable (Protitch 1937) multi-wavelength studies of RW Tri have provided much insight into the system, but have yet to yield an accurate measurement of the masses of the two component stars. By observing the velocities of the component stars, their masses can be calculated using Kepler's laws. A number of authors have attempted to measure the radial velocity of the white dwarf (e.g. Kaitchuck \\ea 1983, Still \\ea 1995 and Mason \\ea 2002), but thus far little work has been done on the secondary star.  Measurement of the secondary star velocity is hampered by the large contrast in optical brightness between it and the accretion disk. Weak  secondary star absorption features were detected in the I-band spectrum of RW Tri by Friend \\ea (1988), and later used by Smith \\ea (1993) to estimate the secondary star velocity as $\\sim250$km/s using a `skew-mapping' technique (cf. Smith \\ea 1998a). Dhillon \\ea (2000) have detected secondary star features in low-resolution K-band spectra of RW~Tri. The relative contribution of the secondary star is higher in the K-band and the equivalent width of the absorption features greater as a consequence.  In this paper we report phase-resolved spectral observations of RW Tri in the K-band, covering the binary orbit. We use these observations to determine the radial velocity of the secondary star in RW Tri. We also re-analyse the original I-band spectra of Smith \\ea (1993).  Based on these results we revisit the mass ratio and the masses of the component stars in RW~Tri and discuss their implications. ",
        "conclusions": "I-band observations of RW Tri yield a secondary star radial velocity amplitude of $250\\pm47$ km/s using the skew mapping technique. K-band observations of RW Tri provide us with a secondary star velocity of $221\\pm29$ km/s which is obtained directly from the observations without using complex mapping methods. The two velocities are consistent within the errors.  We estimate the rotational velocity of the secondary star to be $120\\pm20$km/s using the K-band UKIRT observations. Combining this velocity with the secondary star radial velocity corrected for non-uniform heating derived from the variation in the absorption line strengths, we find a mass ratio range of $0.6-1.1$, which contains the lower adiabatic response limit ($q_{ad,fc}=2/3$). The uncorrected radial velocity results lead to a mass ratio range of $0.5-0.8$.  Combining the radial velocity amplitudes of the secondary star with the radial velocity of the primary determined from  He II emission lines in the optical (Still \\ea 1995) and narrow absorption lines in the UV (Mason \\ea in prep.), yields a range of mass ratios of $1.0-1.3$ and $1.2-1.7$ respectively, based on a range in the secondary heating correction of $\\Delta K=0-24\\%$. The UV mass ratio range lies above the critical mass ratio ($q_{crit}=4/3$); this range is also inconsistent with the rotational velocity results, indicating that the UV velocity very likely includes a non-orbital component.  The He II mass ratio range lies between $q_{crit}=4/3$ and q=5/6, but is only marginally consistent with the measured secondary rotational velocity, and may also contain a non-orbital component. By combining the data on the rotational broadening of the secondary with its measured orbital velocity, with no assumptions regarding the white dwarf velocity, we find most likely values of the primary and secondary masses that lie in the range $0.4-0.7M_{\\odot}$ and $0.3-0.4M_{\\odot}$ respectively, depending on the degree of secondary star heating. The most likely value of $K_1$ is predicted to be 120-130 km/s."
    },
    "0212/astro-ph0212032_arXiv.txt": {
        "abstract": "{We discuss spectral energy distributions of a sample of Herbig Ae/Be stars in the context of a passive irradiated disk model. The data have been presented earlier by \\citet{2001A&A...365..476M}, and preliminary interpretations of these data were given in that paper.  While the spectra of Herbig Ae stars all show similarities, there is significant variation between the spectra, in particular in the shape of the mid-IR rise and in the presence or absence of a silicate feature.  We explore the hypothesis that all these different spectra can be interpreted as pure disk spectra without additional components. Using the model of Dullemond, Dominik and Natta (\\citeyear{2001ApJ...560..957D}) we deduce the disk parameters of a number of the sources, and find that for a large fraction of investigated sources, satisfactory fits can be obtained.  The derived model parameters show that some group Ia sources can only be fit with radially increasing surface densities, indicating the presence of depleted inner disk regions.  The steep-sloped SEDs of group IIa sources can be fit with very compact disks, probably representing disks with collapsed outer regions. The largest difficulties arise from sources that do not show significant silicate emission features.  Our attempts to explain these objects with a pure geometric effect are only partially successful.  It seems that these stars indeed require a strong depletion of small silicate grains.} ",
        "introduction": "\\cdrev{Herbig Ae/Be stars (herafter referred to as HAEBE stars) are thought to be young stellar objects of intermediate mass \\citep{1960ApJS....4..337H,1972ApJ...173..353S,1984A&AS...55..109F,1994A&AS..104..315T,1998A&A...331..211M}. They are associated with large amounts of circumstellar matter, both gas and dust, and are found in or near star forming regions. In the seminal paper of \\citet{1960ApJS....4..337H}, HAEBEs were defined to be near a reflection nebulosity, because this ensured they are young and close to a star forming region.  The close proximity to molecular cloud material in combination with limited angular resolution observations makes it hard to distinguish dust and gas emission from the loose surroundings from that of the circumstellar disk. However, objects have been found that are very similar to other HAEBEs but lack a reflection nebulosity \\citep[see e.g.][]{1998ARA&A..36..233W}.  These stars are called \\emph{isolated Herbig stars}.  Hipparcos observations of the nearest HAEBE stars have confirmed their pre-main-sequence nature \\citep{1997A&A...324L..33V}.}  \\cdrev{HAEBE stars have a strong infrared (IR) excess due to circumstellar dust, often carrying a considerable fraction of the total luminosity of the system. Mid-IR imaging has shown that some HAEBE stars have spatially extended emission at these wavelengths \\citep{1994A&A...292..593P,1994ApJ...432..710D,1998ApJ...509..324D}; these stars often (but not always) are of (early) B spectral type, and their emission can be understood in terms of heating of the surrounding molecular cloud by the strong UV radiation field of the young star, possibly in combination with a disk \\citep{1999ApJ...520L.115M,1998A&A...336..565H}.  However, the less luminous late B, and A type stars in the HAEBE class often show unresolved or compact emission in the mid-IR when observed with 4 meter class telescopes on a scale of 1--2 arcsec, corresponding to 100-200 AU at typical distances of 100\\,pc.  Van den Ancker et al (private communication) have unsuccessfully tried to resolve several of the sources in the sample (HD\\,100546, HD\\,104237, HD\\,144432, HD\\,163296) with the ESO 3.6m telescope, finding only upper limit of typically 0.7''.  Other observations have succeeded resolving the emission of some stars (e.g. AB Aurigae \\citep{1995ApJ...451..777M}, HD100546 \\citep{2001AJ....122.3396G}, HD97048 (van Boeckel et al, in preparation)) and found compact sizes.   The isolated Herbig stars usually show little or no reddening at optical wavelengths, which is remarkable given their large L$_{\\rm IR}$/L$_{*}$ ratio.  An obvious way to explain this observation is to assume a flattened, disk-like distribution of the gas and dust viewed at some intermediate inclination angle.  This disk can be understood as the passively heated remnant of the accretion disk which was present during the initial phase of star formation, and is similar to disks seen around the lower mass T~Tauri stars. Since such passive disks are believed to be the site of planet formation, HAEBE stars have gained considerable attention in recent years.}  \\cdrev{Obviously, the best way to determine the geometry of the circumstellar material around isolated HAEBE stars is by direct imaging.  \\citet{1997ApJ...490..792M} spatially resolved the millimeter continuum and line emission from the Herbig Ae star HD163296 and found an elongated structure. Submillimeter aperture synthesis images in CO lines show rotation profiles in a number of Herbig Ae stars, with rotational velocities consistent with a Keplerian disk \\citep{1998A&A...338L..63D,1999alma.confE..50Q,1997ApJ...490..792M,2000ApJ...529..391M}. Based on this evidence one may conclude with some confidence that at least the outer parts of the circumstellar matter distribution around some (isolated) Herbig Ae stars are, just as in the case of T-Tauri stars, in fact rotating circumstellar disks.}  \\cdrev{Observations of the spatial distribution of the circumstellar matter closer to the star (inwards of about 100 AU) have not yet resulted in a clear picture. There seems to be conflicting evidence suggesting both a flattened and more spherical distribution of gas and dust. Near-IR and optical imaging using ground-based adaptive optics and HST in several cases shows the emission to be flattened \\citep{2000ApJ...544..895G,2001AAS...198.7716G,2001AJ....122.3396G,2000A&A...361L...9P}. However, interferometric observations at 2 $\\mu$m, probing the inner regions on scales below 10 AU, have shown little or no evidence for a disk-like geometry \\citep{2001ApJ...546..358M}.}  In view of the scarcity of resolved observations at these spatial scales, it seems that the procedure of fitting a model to the observed SED is still a valuable tool to study the properties of circumstellar matter around young stars. Under the assumption that the circumstellar dust is in radiative equilibrium, the temperature of the dust can be calculated by solving the radiative transfer equation. By varying the density distribution and the dust composition, one can eventually find a good 'fit' to the data.  One has to keep in mind however that this procedure can not uniquely determine the geometry and density distribution of the circumstellar matter. Often, a multitude of different parameter sets lead to equally good fits \\citep{1994A&A...287..493T}. This clearly means that the SED alone does not contain enough information about the geometry and dust distribution in these objects.  Nevertheless, model fitting can retrieve interesting information from the SED if one uses additional constraints to the density distribution. These constraints can be obtained from e.g.~spectroscopic quantities (lines, features), or any kind of imaging data when available. In addition to this, one can add {\\em physics} to the model. For instance, one could require the density distribution to be in hydrostatic equilibrium: a self-consistent disk model. A simple, but powerful model of this kind is the semi-analytic passive flaring disk model of \\citet{1997ApJ...490..368C}. This model naturally reproduces the shape of the spectrum of many Herbig Ae stars longwards of 6 micron. Moreover, as was first noted by \\citet{2001A&A...371..186N}, if one correctly accounts for emission from the disk's inner rim, also the SED shortwards of 6 microns (the conspicuous 3 micron bump) can be explained. In a recent paper, Dullemond, Dominik \\& Natta (2001, henceforth DDN), presented such a self-consistent disk model including both the emission from the flaring part of the disk and from the inner rim. They found that emission from the hot inner rim (being hotter than the flaring disk and thus having a higher scale height) can account for the near-IR excess in the HAEBE star AB~Aur; the DDN model for AB~Aur is the first disk-only model that can account for the \\emph{entire} SED of AB~Aur.  Encouraged by our success in fitting AB~Aur, we decided to investigage the applicability of the DDN model to a wider sample of HAEBE stars. \\cdrev{We use the sample described by \\citet{2001A&A...365..476M}. This is a set of isolated Herbig Ae stars which was selected from a larger set of Herbig AeBe stars \\citep{1994nesh.conf..405W,1994PhDT.......226B,1998A&A...331..211M}. To our knowledge, none of the stars (with the exception of AB~Aur) shows any significant reflection nebulosity or significant extended emission on scales much larger than expected for a disk.  This suggests that the SED is not contaminated by large scale dust clouds surrounding the object.  The objects were further selected by their very low circumstellar extinction - typically 0.5 magnitudes or less.  Exceptions are HD\\,142666 (0.9\\,mag) and HD\\,150193 (1.5\\,mag) \\citep[and private communication]{1998A&A...330..145V}}.  In their paper, Meeus et al.~presented combined ISO SWS/LWS spectra and literature photometry data of 14 Herbig Ae/Be stars, and found that these objects show both striking similarities in the spectrum, and important differences.  We will attempt to interpret this variety of spectra entirely in the context of the DDN model, without additional components. This way we can find out whether an isolated disk picture is sufficient to explain the SEDs of these stars. And if we can, we will see what we can learn from the deduced parameters.  The paper is organized as follows.  In section \\ref{sec:sample-seds} we describe the sample.  In section~\\ref{sec:model-fitt-proc} the DDN model is introduced and we detail the fitting procedure.  In section~\\ref{sec:results}, the obtained fits are shown and discussed.  Finally, in section~\\ref{sec:discussion}, we reflect on the successes and failures of the model to reproduce the SEDs of isolated Herbig stars. ",
        "conclusions": ""
    },
    "0212/astro-ph0212518_arXiv.txt": {
        "abstract": "We consider a model of interacting cosmological constant/quintessence, where dark matter and dark energy behave as, respectively, two coexisting phases of a fluid, a thermally excited Bose component and a condensate, respectively. In a simple phenomenological model for the dark components interaction we find that their energy density evolution is strongly coupled during the universe evolution. This feature provides a possible way out for the coincidence problem affecting many quintessence models. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/astro-ph0212342_arXiv.txt": {
        "abstract": "EUSO experiment, while monitoring the downward Earth atmosphere layers, may observe among common Ultra High Energy Cosmic Rays, UHECR, also High Energy Neutrino-Induced Showers either blazing upward to the detectors at  high ($\\sim$ PeVs) energies or at much higher GZK, $\\sim E_{\\nu}\\geq 10^{19}$ eV energies, showering horizontally in air or vertically downward.  A small fraction of these upward, horizontal and vertical Shower maybe originated by direct astrophysical UHE neutrino interacting on terrestrial air layers itself; however the dominant UHE neutrino signal are Upward and Horizontal Tau Air-Showers, UPTAUS and HORTAUs (or Earth skimming $\\nu$), born within widest Earth Crust Crown (Sea or Rock) Areas, by UHE $\\nu_{\\tau} + Nuclei$ $\\rightarrow \\tau$ interactions, respectively at PeVs and GZK energies: their rate and signatures are shown in a neutrino fluence map for EUSO thresholds versus other UHE air interacting neutrino signals and backgrounds. The effective target Masses originating HORTAUs seen by EUSO may exceed (on sea) a wide and huge ring volume $\\simeq 5130$ $km^3$. The consequent HORTAUS event rate (even at $10\\%$ EUSO duty cycle lifetime) may deeply test the expected Z-Burst models by at least a hundred of yearly events. Even rarest but inescapable GZK neutrinos (secondary of photopion production of observed cosmic UHECR) might be discovered in a few (or a tens) horizontal shower events; in this view an extension of EUSO detectability up to $\\sim E_{\\nu}\\geq 10^{19}$eV threshold is to be preferred. A wider collecting EUSO telescope (3m diameter) might be considered. ",
        "introduction": "While longest ${\\mu}$ tracks in $km^3$ underground detector have been, in last three decades, the main searched UHE neutrino signal, Tau Air-showers by UHE neutrinos generated in Mountain Chains or within Earth skin crust at Pevs up to GZK ($>10^{19}$ eV) energies have been recently proved to be a  new powerful amplifier in Neutrino Astronomy \\cite{Fargion et all 1999}, \\cite{Fargion 2000-2002},\\cite{Bertou et all 2002},\\cite{Hou Huang 2002},\\cite{Feng et al 2002}. This new Neutrino $\\tau$ detector will be (at least) complementary to present and future, lower energy, $\\nu$ underground  $km^3$ telescope projects (from AMANDA,Baikal, ANTARES, NESTOR, NEMO, IceCube). In particular Horizontal Tau Air shower may be naturally originated by UHE $\\nu_{\\tau}$ at GZK energies crossing the thin Earth Crust at the Horizon showering far and high in the atmosphere \\cite{Fargion 2000-2002},\\cite{Fargion2001a}, \\cite{Fargion2001b},\\cite{Bertou et all 2002},\\cite{Feng et al 2002}. UHE $\\nu_{\\tau}$ are abundantly produced by flavour oscillation and mixing from muon (or electron) neutrinos, because of the large galactic and cosmic distances respect to the neutrino  oscillation ones (for already known neutrino mass splitting). Therefore EUSO may observe many of the above behaviours and it may constrains among models and fluxes and it may also answer open standing questions. I will briefly enlist, in this first preliminary presentation, the main different signatures and rates of UHECR versus UHE $\\nu$ shower observable by EUSO at 10\\% duty cycle time within a 3 year record period, offering a first estimate of their signals. Part of the results on UHECR are probably  well known, even here it is re-estimated. Part of the results, regarding the UPTAUs and HORTAUs, are new and they rule the UHE $\\nu$ Astronomy in EUSO. \\begin{figure}\\centering\\includegraphics[width=8cm]{Fig001HIAS.eps} \\vspace{-1.5cm} \\caption {A very schematic Horizontal High Altitude Shower (HIAS); its fan-like imprint is due to geo-magnetic bending of charged particles at high quota ($\\sim 44 km$). The Shower may  point to an satellite as old gamma GRO-BATSE detectors or very recent Beppo-Sax,Integral, HETE, Chandra or future Agile and Swift ones. \\cite{Fargion 2000-2002},\\cite{Fargion2001a},\\cite{Fargion2001b}. The HIAS Showers is open and forked in  five (or three or at least two main component): ($e^+,e^-,\\mu^+,\\mu^-, \\gamma $, or just positive-negative); these  multi-finger tails may be seen as split tails  by EUSO. } \\label{fig:fig1} \\vspace{-0.3cm} \\end{figure} \\begin{figure} \\vspace{- 0.2cm} \\centering\\includegraphics[width=8cm]{Fig002HORTAUS.eps} \\vspace{-1cm} \\caption { As above  Horizontal Upward Tau Air-Shower (HORTAUS) originated by UHE neutrino skimming the Earth: fan-like jets due to geo-magnetic bending  shower at high quota ($\\sim 23-40 km$): they may be pointing to an orbital satellite detector \\cite{Fargion 2000-2002}, \\cite{Fargion2001a}, \\cite{Fargion2001b}. The Shower tails may be also observable by EUSO just above it.} \\label{fig:fig2} \\end{figure} \\begin{figure} \\vspace{-0.5cm} \\centering\\includegraphics[width=8cm]{Fig003EusoUPTAUS.eps} \\caption {A very schematic Upward Tau Air-Shower (UPTAUs)  and its open fan-like jets due to geo-magnetic bending at high quota ($\\sim 20-30 km$). The gamma Shower may be pointing to an orbital detector \\cite{Fargion 2000-2002}, \\cite{Fargion2001a}, \\cite{Fargion2001b}. Its vertical Shower tail may be spread by geo-magnetic field into a thin eight-shape beam observable  by EUSO  as a small blazing oval (few dot-pixels) aligned orthogonal to the local magnetic field .} \\label{fig:fig3} \\vspace{-0.9cm} \\end{figure} ",
        "conclusions": "Highest Energy Neutrino signals may be well observable by next generation satellite as EUSO: the main source of such neutrino traces are UPTAUs (Upward Tau blazing the telescope born in Earth Crust) and mainly HORTAUs (Horizontal Tau Air-Showers originated by an Earth-Skimming UHE $\\nu_{\\tau}$). These showers will be opened in a characteristic thin fan-jet ovals like the $8$-shape horizontal cosmic ray observed on Earth. The UPTAUs will arise mainly at PeV energies (because the Earth neutrino opacity at higher energies and because the shorter $\\tau$ boosted lenght , at lower energies)\\cite{Fargion 2000-2002}; UPTAUs will be detected as a thin stretched multi-pixel event by EUSO, whose orientation is polarized orthogonal to local geo-magnetic field. The EUSO sensibility (effective volume ($V_{eff}$$\\sim 0.1 km^3$) for 3 years of detection) will be deeper an order of magnitude below present AMANDA-Baikal bounds. Horizontal Tau Air-Shower at GZK energies will be better searched and revealed. They are originated along huge Volumes around the EUSO Area. Their horizontal skimming secondary $\\tau$ decay occur far away $\\geq 500$ km, at high altitude ($\\geq 20-40$ km) and it will give clear signals distinguished from downward horizontal UHECR. HORTAUs are grown by UHE neutrino interactions inside huge volumes ($V_{eff}$$\\geq 5130-6250 km^3$) respectively for incoming neutrino energy $E_{\\nu_{\\tau}}$ $\\simeq 10^{19}$ eV and $3 \\cdot10^{19}$ eV. To obtain these results we applied the procedure described in recent articles \\cite{Fargion 2002b},\\cite{Fargion 2002c},\\cite{Fargion 2002d}. As summirized in last Figures the expected UHE fluence $$\\Phi_{\\nu}\\simeq 10^{3} eV cm^{-2} s^{-1}$$ needed in most Z-Shower models (as well as in most topological relic scenario) to solve GZK puzzles, will lead to nearly a hundred of horizontal events a year comparable to UHECR ones. Even in the most conservative scenario where a minimal GZK-$\\nu$ fluence must take place (at least at $$\\Phi_{\\nu} \\simeq 10 eV cm^{-2} s^{-1}$$ ,just comparable to well observed Cosmic Ray fluence), a few or a ten of such UHE astrophysical neutrino must be observed (respectively at $10^{19}$ eV and $3 \\cdot10^{19}$ eV energy windows) during three year of EUSO data recording. To improve their visibility EUSO must, in our opinion one may:\\\\ a) Improve the fast pattern recognition of Horizontal Shower Tracks with their few distant dots with forking signature.\\\\ b) Enlarge the Telescope Radius to embrace also lower $10^{19}$ eV energy thresholds where UHE neutrino signals are enhanced.\\\\ c) Consider a detection  at  angular $\\Delta\\theta$ and at height $\\Delta h$ level within an accuracy $\\Delta\\theta \\leq 0.2^o$,$\\Delta h \\leq 2$ km.\\\\ Even all  the above results have been derived carefully  following \\cite{Fargion 2002b},\\cite{Fargion 2002c},\\cite{Fargion 2002d} in a minimal realistic framework they may be used within $10\\%$ nominal value due to the present uncertain in  EUSO detection capabilities.         \\subsection*{Acknowledgment} The author wish to thank Prof. Livio Scarsi  for inspiring the present search  as well as the EUSO collaboration for the exciting discussion during  November workshop in Rome; the author thanks also C.Leto and P.G.De Sanctis Lucentini and M.Teshima for support and technical suggestions."
    },
    "0212/astro-ph0212497_arXiv.txt": {
        "abstract": "We constrain the basic comological parameters using the first observations by the Very Small Array (VSA) in its extended configuration, together with existing cosmic microwave background data and other cosmological observations. We estimate cosmological parameters for four different models of increasing complexity. In each case, careful consideration is given to implied priors and the Bayesian evidence is calculated in order to perform model selection. We find that the data are most convincingly explained by a simple flat $\\Lambda\\rm{CDM}$ cosmology without tensor modes. In this case, combining just the VSA and COBE data sets yields the 68 per cent confidence intervals $\\Omega_{\\rm{b}}h^2=0.034 ^{+0.007} _{-0.007}$, $\\Omega_{\\rm{dm}}h^2=0.18 ^{+0.06}_{-0.04}$, $h=0.72^{+0.15}_{-0.13} $, $n_s=1.07 ^{+0.06 }_{-0.06}$ and $\\sigma_8=1.17 ^{ +0.25 }_{ -0.20}$.  The most general model considered includes spatial curvature, tensor modes, massive neutrinos and a parameterised equation of state for the dark energy. In this case, by combining all recent cosmological data, we find, in particular, 95 percent limit on the tensor-to-scalar ratio $R < 0.63$ and on the fraction of massive neutrinos $f_\\nu < 0.11$; we also obtain the 68 per cent confidence interval $w=-1.06^{+0.20}_{-0.25}$ on the equation of state of dark energy. ",
        "introduction": "In the past two years, a number of experiments have produced accurate measurements of the power spectrum of anisotropies in the cosmic microwave background (CMB) radiation on a range of angular scales (Hanany et al. 2002; Netterfield et al. 2002; Halverson et al. 2002; Sievers et al. 2002; Benoit et al. 2002). These data, together with other cosmological observations, have been used to place increasingly tight constraints on the values of cosmological parameters in current models of the formation and evolution of structure in the Universe.  In this letter, we repeat this process with the inclusion of the latest observations from the Very Small Array (VSA) in its extended configuration. Results from the VSA in its compact configuration have already been presented in Watson et al. (2002), Taylor et al. (2002), Scott et al. (2002) and Rubi\\~no-Martin et al. (2002) (hereafter Papers I - IV). In Grainge et al. (2002, hereafter Paper V) these data are combined with the new extended configuration observations to produce a combined power spectrum with 16 spectral bins spanning the range $\\ell=160-1400$.  This joint set of observed band-powers provides powerful new constraints on cosmological parameters. In this letter we extend the traditional likelihood approach of previous analyses to a fully Bayesian treatment, including careful consideration of our knowledge of cosmological parameters prior to the inclusion of any data, and the calculation of Bayesian evidences to perform model comparisons.  \\begin{table} \\begin{center} \\caption{The priors assumed for the basic parameters common to all four cosmological models under consideration. The notation $(a,b)$ for parameter $x$ denotes a top-hat prior in the range $a<x<b$} \\label{prior} \\begin{tabular}{lc} \\hline\\hline Basic parameter & Prior \\\\ \\hline $\\omega_b$ & $(0.005,0.80)$ \\\\ $\\omega_{dm}$ & $(0.01,0.9)$ \\\\ $h$ &  $(0.4,1.0)$ \\\\ $n_s$ &  $(0.5,1.5)$ \\\\ $z_{re}$ & $(4,20)$\\\\ $10^{10} A_s$ & $(0,100)$ \\\\ \\hline\\hline \\end{tabular} \\end{center} \\end{table} ",
        "conclusions": "By comparing the effective priors on the cosmological parameters plotted in Fig.~\\ref{graph-implied} with the posterior distributions inferred from the VSA and COBE data alone (Fig.~\\ref{vsae-only}), we see that the main strength of the VSA lies in its large $\\ell$-range, which allows one to constrain a wide variety of cosmological parameters. In particular, we observe that the VSA data significantly improves the constraints on $\\omega_{\\rm b}$, $\\omega_{\\rm dm}$, $\\Omega_{\\rm k}$, $\\sigma_8$ and $n_s$.  We note that, as a direct consequence of the pronounced third peak in the VSA power spectrum, the preferred value for $\\omega_{\\rm b}$ is somewhat larger than that from the nucleosynthesis constraint (Burles et al. 2001) or that from the combined CMB data. This excess is statistically significant only at the $1.6 \\sigma$ level. VSA and COBE data alone do not, however, provide significant new constraints on either $h$ or $\\Omega_\\Lambda$, especially for Model B.  On comparing Figs~\\ref{vsae-only} and \\ref{cmb-only}, we see that, as one would expect, the constraints on all parameters are tightened by the inclusion of all the CMB data. Indeed, for Model A, reasonable constraints are obtained on all the parameters under consideration, except for $z_{\\rm re}$, although there is some indication that low values of $z_{\\rm re}$ are preferred. For many of the parameters, the constraints are not significantly broader when one allows the spatial curvature to vary in Model B. Nevertheless, for this model, some parameters do become relatively unconstrained, in particular $h$, $\\Omega_{\\rm m}$ and $\\Omega_\\Lambda$, and the limits on the age of the Universe are widened considerably.  Using all recent available cosmological data sets allows one to place tight constraints on nearly all parameters, even for Model B. Indeed, only $z_{\\rm re}$ remains relatively unconstrained, but the earlier indication that lower values are preferred appears to be reinforced. We note that, aside from the parameters $h$ and the age of the Universe, the constraints obtained are very similar for Model A and Model B. In fact, from Table~\\ref{results}, we see that, even for Models C and D, the constraints on the parameters plotted in Fig.~\\ref{all-only} are not significantly broadened. In Fig.~\\ref{extrapars}, we also plot the constraints obtained on the additional parameters $R$, $n_t$, $f_\\nu$ and $w$ for Model D. We see that the data favour low values of $R$ and high values of $n_t$. More significant constraints can be placed on $f_\\nu$ and $w$. We note from the figure that the marginalised distribution for each parameter possesses a single peak, although these are not pronounced features, especially for $f_\\nu$. Of particular interest is that the preferred value of $w$ corresponds to a cosmological constant and that a firm upper limit is obtained on the fraction of dark matter in the form of massive neutrinos.  To distinguish between the different cosmological models considered here, we have calculated the Bayesian evidence for each model-data set combination. It must be remembered that evidence values can only be compared between different models using the same data set.  From Table~\\ref{results}, we see that, for each data set, Model A (i.e. the simple 6-parameter $\\Lambda$CDM cosmology) is preferred, although the evidence ratio between Model A and Model B is of order unity in each case. For Models C and D, however, the evidence ratio with Model A is considerably larger and shows clearly that such general models are not necessary to explain current cosmological observations.   Since Model A is found to have the largest evidence, it is of interest to consider more closely the impressively tight parameter constraints that can be achieved in this simple case using all the available cosmological data. In particular, let us focus on $h$ and $\\omega_{\\rm b}$, which are parameters directly probed by the HST Key Project (Freedman et al. 2001) and nucleosynthesis data (Burles et al. 2001). In Fig.~\\ref{hbbn}, we plot the posterior distributions of these parameters for Model A, with (thick line) and without (thin line) the corresponding direct constraint (dashed line). In each case, we see that the direct constraint is overwhelmed by the combination of the other data sets.  It is clear from the above analyses that the constraints on the parameters in the simplest models become impressively small when a wide range of cosmological probes are used. It remains to be seen, however, whether this level of precision is matched by a corresponding level of accuracy in the values determined. The model comparison exercise demonstrated here does not address the possibility that different experiments have different systematic errors. Nevertheless, one would hope that these might be reduced when many experiments are combined. In future work, we will investigate other models, with different combinations of free and fixed parameters and perform a hyperparameter analysis (Hobson et al. 2002) to reveal any discrepancies between data sets."
    },
    "0212/astro-ph0212174_arXiv.txt": {
        "abstract": "We present near-infrared magnitudes for all white dwarfs (selected from the catalog of McCook \\& Sion) contained in the 2 Micron All Sky Survey Second Incremental Data Release. We show that the near-IR color-color diagram is an effective means of identifying candidate binary stars containing a WD and a low mass main sequence star.  The loci of single WDs and WD + red dwarf binaries occupy distinct regions of the near-IR color-color diagram.  We recovered all known unresolved WD + red dwarf binaries located in the 2IDR sky coverage, and also identified as many new candidate binaries (47 new candidates out of 95 total). Using observational near-IR data for WDs and M--L dwarfs, we have compared a sample of simulated WD + red dwarf binaries with our 2MASS data.  The colors of the simulated binaries are dominated by the low mass companion through the late-M to early-L spectral types.  As the spectral type of the companion becomes progressively later, however, the colors of unresolved binaries become progressively bluer.  Binaries containing the lowest mass companions will be difficult to distinguish from single WDs solely on the basis of their near-IR colors. ",
        "introduction": "In the search for extrasolar planets, various methods have been employed to detect the signatures of faint stellar and sub-stellar companions.  For main sequence primary stars, faint low mass companions are often hidden in the glare of the more luminous primary, and radial velocity variations are small and therefore difficult to detect.  On the other hand, observing low mass companions to white dwarfs (WDs) offers many advantages compared to main sequence primaries.  Since WDs are less luminous than main sequence stars, the brightness contrast compared to a potential faint companion is significantly reduced. Most importantly, the markedly different spectral energy distributions of the WDs and their low mass companions makes the detection and separation of the two components relatively straightforward even with simple broad-band multi-color photometry.  Because WDs have traditionally been identified and studied via observations in the blue part of the spectrum, comparatively little is known about their infrared (IR) properties.  The recent discovery that very cool WDs are much bluer in the IR than previously thought to be the case \\citep{hodgkin00} highlights how little is known about WD spectral properties at longer wavelengths.  Consequently, we are studying the group near-IR photometric properties of WDs. In this paper, we present analysis of the near-IR color-color diagram of WDs, which demonstrates a means of identifying candidates for WDs with close (unresolved), cool, low mass stellar or sub-stellar companions. ",
        "conclusions": "We have shown that the near-IR color-color diagram is an effective means of identifying candidate binary stars containing a WD and a low mass main sequence star.  The loci of single WDs and WD + red dwarf binaries occupy distinct regions of the near-IR color-color diagram.  Using our data from the 2MASS 2IDR, we recovered all known unresolved WD + red dwarf binaries located in the 2IDR sky coverage, and also identified nearly as many new candidate binaries (47 new candidates out of 95 total).  In addition, a handful of the known resolved binaries may actually be triple systems, in which the WD component is itself an unresolved WD + red dwarf binary.  We expect to be able to more than double again the number of candidate binaries using the forthcoming full sky data release from 2MASS.  Using observational near-IR data for WDs and M--L dwarfs, we have compared a sample of simulated WD + red dwarf binaries with our 2MASS data.  The colors of the simulated binaries are dominated by the low mass companion through the late-M to early-L spectral types.  As the spectral type of the companion becomes progressively later, however, the colors of unresolved binaries become progressively bluer.  Binaries containing the lowest mass companions will be difficult to distinguish from the locus of single WDs in the near-IR color-color diagram.  We have identified an additional 15 WDs that may comprise such binaries.  It is encouraging that one of these has been found to be a close WD + red dwarf binary in which the companion has a mass of $\\lesssim0.1M_{\\odot}$ \\citep{marsh96}.  In order to distinguish the WD + red dwarf binaries from single WDs for systems containing the lowest mass L dwarfs (and brown dwarfs), it is likely to be necessary to observe further into the infrared; for example, at the mid-IR wavelengths observable with the {\\em Space Infrared Telescope Facility} (e.g., \\citealt{ignace02})."
    },
    "0212/astro-ph0212160_arXiv.txt": {
        "abstract": "The importance of studying old elliptical galaxies at redshift $z \\sim 1.5$ is reviewed, considering both what can be learned by extending studies of the evolution of cluster galaxy scaling relations to earlier cosmic epochs, and the age--dating of old elliptical galaxies at high redshifts. Following this, the first results are provided of an on--going project to find such distant elliptical galaxies, through an investigation of the cluster environments of powerful radio sources with redshifts $1.44 < z < 1.7$. These studies show a considerable excess of red galaxies in the radio sources fields, with the magnitudes ($K \\gta 17.5$) and colours ($R-K > 4$) expected of old passively evolving galaxies at the radio source redshift. The red galaxy overdensities are found on two different scales around the radio sources; a pronounced small--scale peak at radial distances of $\\lta 150$\\,kpc, and a weaker large--scale excess extending out to 1 - 1.5\\,Mpc. The presence and richness of these red galaxy excesses varies considerably from source--to--source. An interpretation of these results is provided. ",
        "introduction": "One of the key questions facing astrophysics today concerns the epoch and mode of formation of the most massive elliptical galaxies. Whilst the majority of galaxies can be easily accommodated by a slow assembly of galaxies from bottom-up, within the standard predictions of semi--analytical models of a cold dark matter Universe \\cite{kau98,col00}, some lines of evidence suggest that an earlier and more dramatic formation mechanism may be necessary to accommodate (some of) the most massive galaxies. In particular:\\\\ \\noindent {\\it 1. Powerful AGN at high redshifts:} the host galaxies of powerful radio sources are the most massive galaxies known at early cosmic epochs [e.g. 3], with stellar masses of a few times $10^{11} M_{\\odot}$ at redshifts $1 \\lta z \\lta 2$ \\cite{bes98}. At higher redshifts, even if fueled at the Eddington limit, central supermassive black holes of $10^{8-9} M_{\\odot}$ are required to produce the observed radio luminosity, suggesting that massive galaxies are already in place at redshifts $z \\gta 5$ \\cite{bre99}.\\\\ \\noindent {\\it 2. The nature of sub-millimetre selected galaxies:} many of the powerful sub-mm selected galaxies seen with SCUBA are objects at $z > 2$, which are currently forming stars at hundreds or thousands of solar masses per year [e.g. 6]\\nocite{sma02}. The K$-z$ relation for these SCUBA galaxies indicates that these are amongst the most massive objects at their epoch \\cite{dun02}. SCUBA galaxies therefore appear to be massive galaxies which are forming the bulk of their stars in a single dramatic starburst at high redshift.\\\\ \\noindent {\\it 3. Cluster galaxy scaling relations:} the very low scatter found around the colour magnitude relation [e.g. 8]\\nocite{ter02} and the fundamental plane (a tight correlation between the surface brightness ($I_{\\rm e}$), characteristic radius ($r_{\\rm e}$), and velocity dispersion ($\\sigma$) of cluster ellipticals, which can be roughly approximated as: $r_{\\rm e} \\propto \\sigma^{1.2} I_{\\rm e}^{-0.8}$; e.g. 9)\\nocite{jor96} of elliptical galaxies in low redshift clusters imply a fairly high formation redshift of these ellipticals, with little bluening from later starbursting activity. The fundamental plane essentially measures the mass--to--light ($M/L$) ratio of the galaxies (from $M \\propto \\sigma^2 r_{\\rm e}$ and $L \\propto I_{\\rm e} r_{\\rm e}^2$ it can be written as $M/ L \\propto M^{\\approx 0.25}$); by comparing clusters at different redshifts [e.g. 10,11],\\nocite{kel97,dok01} the mass--to--light ratio is found to evolve very slowly with redshift out to $z \\sim 0.8$, as $\\Delta Log(M/L) \\propto 0.4z$.  Further, the scatter around neither the fundamental plane nor the colour--magnitude relation shows any significant increase out to redshifts $z \\sim 1$ \\cite{dok01,sta98}. These results both further support the argument that the elliptical galaxies must have formed at a very early cosmic epoch, with limited later star formation (although see \\cite{dok01} for discussion of a possible `progenitor--bias' selection effect).\\\\ \\noindent {\\it 4. The ages of massive ellipticals at high redshift:} Dunlop et~al.  \\cite{dun96} age--dated the old red elliptical galaxy 53W091 at $z=1.55$ through a detailed investigation of its optical spectral energy distribution. They derived an age of 3.5\\,Gyr, comparable to the age of the Universe at that redshift, again implying a very early formation epoch. Although there has since been considerable debate as to the age of this galaxy, and of the similar 53W069, Nolan et~al \\cite{nol01} convincingly defend this initial age estimate. Similar results are beginning to be obtained for Extremely Red Objects (EROs; e.g \\cite{cim02} and refs therein).  The identification of a sample of cluster elliptical galaxies at redshifts $z \\sim 1.5$ would provide a key opportunity to advance our understanding of the epoch and mode of formation of these objects. This is an early enough epoch to test the `progenitor--bias' model (which would allow later star formation in cluster ellipticals), and would also offer a unique opportunity to compare, through age--dating techniques, the ages of the most luminous elliptical galaxies in the cluster with those of less luminous cluster ellipticals, and therefore directly observe whether there is an age--luminosity relation for cluster ellipticals. Redshift 1.5 is a critical cosmic epoch for these studies: it is still sufficiently nearby that cluster ellipticals can still be identified on the basis of multi--colour imaging in a practical observing time, and is also an epoch where elliptical galaxies are still expected to be plentiful. In this article, the first steps to providing such a sample of elliptical galaxies are discussed. ",
        "conclusions": ""
    },
    "0212/astro-ph0212356_arXiv.txt": {
        "abstract": "We investigate galaxy clustering and the correlations between galaxies and mass in the $\\Lambda$CDM cosmological model (inflationary cold dark matter with $\\Omega_m=0.4$, $\\Omega_\\Lambda=0.6$, $h=0.65$, $n=0.95$, $\\sigma_8=0.8$), using a large, smoothed particle hydrodynamics simulation (SPH, with $2\\times 144^3$ particles in a $50\\hmpc$ cube).  Simulated galaxies can be unambiguously identified as clumps of stars and cold gas a few kpc to a few tens of kpc across, residing in extended halos of hot gas and dark matter; the space density of the resolved galaxy population at $z=0$ corresponds to that of observed galaxies with luminosity $L \\ga L_*/4$.  We investigate the galaxy correlation function, the pairwise velocity dispersion and mean pairwise velocity, and the second and third moments of counts-in-cells; we also investigate the galaxy-mass correlation function and the average extended mass distributions around galaxies, both of which can be measured via galaxy-galaxy lensing.  For the most part, the predicted biases between galaxies and dark matter lead to good agreement with current observations, including: (1) a nearly constant comoving correlation length from $z=3$ to $z=0$ for mass-selected galaxy samples of constant comoving space density; (2) an rms bias factor $b_\\sigma \\approx 1$ at $z=0$; (3) a scale-dependent bias on small scales that transforms the curved dark matter correlation function into a nearly power-law galaxy correlation function; (4) galaxy pairwise dispersion and hierarchical skewness ratio $S_3$ in good agreement with observed values, and lower than values for the dark matter by $\\sim 20\\%$; (5) a ratio of galaxy-galaxy to galaxy-mass correlation functions consistent with recent measurements from the Red Cluster Sequence survey; and (6) a mean excess mass $\\Delta M(260\\hkpc)$ approximately proportional to galaxy baryon mass $M_b$, in agreement with estimates from the Sloan Digital Sky Survey (SDSS).  All of these clustering properties vary with galaxy baryon mass and, more strongly, with the age of a galaxy's stellar population.  The predicted dependences are in good qualitative agreement with the observed dependence of galaxy clustering and the galaxy-mass correlation function on galaxy type.  The predicted ratio $\\Delta M(260\\hkpc)/M_b$ is lower than the SDSS estimates by a factor of $\\sim 1.5-2$ for galaxies with $M_b \\ga 2\\times 10^{11}M_\\odot$.  A test with a higher resolution (smaller volume) simulation suggests that this discrepancy is mostly a numerical artifact; if so, then the SDSS weak lensing comparison leaves little room for feedback or other astrophysical processes to reduce the stellar masses of luminous galaxies, at least given our adopted cosmological parameters.  On the whole, our results show that the $\\Lambda$CDM model and the galaxy formation physics incorporated in the SPH simulation give a good account of observed galaxy clustering, but anticipated improvements in clustering and weak lensing measurements will soon test this picture in much greater detail. ",
        "introduction": "\\label{sec: intro}  The clustering of galaxies has long been an essential testing ground for cosmological models and the theory of galaxy formation, with comparisons between predicted and observed clustering driving much of the progress in cosmology during the 1970s and 1980s.  Such comparisons provided early support for gravitational instability as the central mechanism of structure formation \\citep[e.g.,][]{davis77}, for initial conditions with a power spectrum redder than white noise \\citep[e.g.,][]{gott75,gott77}, for a ``bottom up'' rather than ``top down'' sequence of gravitational clustering \\citep[e.g.,][]{white83,davis85,fry85}, for an approximately Gaussian distribution of primordial fluctuations \\citep[e.g.,][]{gott89,weinberg92,bouchet93}, and for a power spectrum significantly redder than that predicted by the ``standard'' ($\\Omega_m=1$, $h\\equiv H_0/100\\kmsmpc=0.5$) cold dark matter scenario \\citep[e.g.,][]{efstathiou90,maddox90,park94}. The growth of galaxy redshift surveys has led to measurements of increasing precision and detail, and cosmological N-body simulations have developed into a powerful tool for calculating the gravitational clustering of collisionless dark matter from specified initial conditions.  The main obstacle to drawing stronger inferences from the data is the dependence of theoretical predictions on {\\it both} the relatively straightforward physics of gravitational clustering, which largely determines the distribution of dark matter, and the more complex physics of galaxy formation, which determines the relation between galaxies and mass, often referred to generically as ``bias.''  In this paper, we study galaxy clustering and galaxy bias in a large ($2\\times 144^3$ particles, $50\\hmpc$ cube), smoothed particle hydrodynamics (SPH) simulation of a low density, inflationary cold dark matter universe with a cosmological constant (\\lcdm).  Our goals are, first, to see whether the currently leading cosmological model and ``standard'' galaxy formation physics as incarnated in our SPH simulation can reproduce existing observations, and, second, to give guidance for the physical interpretation of those observations and for the design of more demanding tests of the theoretical predictions. The N-body+hydro approach to {\\it ab initio} predictions of galaxy bias has a substantial history, beginning with dissipative ``sticky particle'' simulations by \\cite{carlberg89} and continuing in the early 1990s with simulations using SPH \\citep{katz92,evrard94} or Eulerian grid hydrodynamics \\citep{cen92} to model galaxy formation in a representative (but small) cosmological volume. The present paper is a direct descendant of Katz et al.\\ (1992), but we study a different cosmological model, we incorporate improvements to the physical treatment of radiative cooling and star formation \\citep[][hereafter KWH]{katz96}, and, above all, we take advantage of the parallel implementation of TreeSPH \\citep{dave97} and advances in computer technology to simulate a volume more than 90 times larger than that of Katz et al.\\ (1992), at similar resolution.  Our analysis overlaps significantly with other recent studies using large volume SPH \\citep{pearce99,pearce01,yoshikawa01} or Eulerian hydrodynamics \\citep{cen00} simulations.  Relative to these investigations, we have higher mass and/or spatial resolution and a somewhat smaller simulation volume, as discussed in \\S\\ref{sec: methods} below.  Two other approaches to {\\it ab initio} predictions of bias have gained currency in recent years: high-resolution, collisionless N-body simulations that identify galaxies with ``subhalos'' in the dark matter distribution \\citep[e.g.,][]{colin99,kravtsov99}, and a hybrid method that combines N-body simulations of the dark matter component with semi-analytic treatments of the galaxy formation physics \\citep[e.g.,][]{kauffmann97,governato98,kauffmann99a,kauffmann99b,benson00a, benson00b,benson01,hatton02}. We will discuss comparisons of our results to those from other hydrodynamic simulations and from the high-resolution N-body and hybrid approaches as they arise.  We carry out a detailed comparison between our simulation and the semi-analytic model of \\cite{benson00a} in a separate paper that focuses on the ``halo occupation distribution'' (HOD) predicted by the two methods \\citep{berlind02b}. The HOD description can be used to calculate many different galaxy clustering statistics, and it helps explain the origin of bias in a physically intuitive manner (see \\citealt{berlind02} and references therein). Based on this comparison, we conclude that our SPH approach and Benson et al.'s (\\citeyear{benson00a}) semi-analytic method should yield similar predictions for most galaxy clustering statistics.  Here we focus mainly on the ``classic'' measurements of galaxy clustering --- the two-point correlation function, variance of galaxy counts, and pairwise velocity moments --- and on one of the simplest measures of higher-order clustering, the third moment of counts-in-cells. Large redshift surveys like the 2dF Galaxy Redshift Survey (2dFGRS) and the Sloan Digital Sky Survey (SDSS) now allow precise measurements of these quantities for multiple classes of galaxies, defined by luminosity, color, morphology, or spectral type (see, e.g., \\citealt{norberg01,norberg02b,zehavi02}, and references therein). While we cannot model these variations with galaxy type in detail --- our simulation volume is far smaller than these surveys, and we do not resolve the morphology of the simulated galaxies --- we can examine the predicted trends of clustering with baryon mass, stellar population age, and local environment.  We concentrate on present-day clustering, but we also compute the evolution of the two-point function, which can be compared to results from deep redshift surveys and to studies of Lyman-break galaxies at $z\\approx 3$ \\citep[e.g.,][]{adelberger98,adelberger02}. The discovery that the clustering of Lyman-break galaxies is similar to that of $L_*$ galaxies at $z=0$, despite the weaker expected clustering of the underlying mass distribution, provides strong evidence that the bright galaxy population was highly biased at $z\\approx 3$, even if galaxies roughly trace mass today.  Galaxy-galaxy weak lensing is an important new observational probe of the relation between galaxies and dark matter, measuring the galaxy-mass cross-correlation function and the extended mass distributions around galaxies of different types and in different environments. We will devote considerable attention to modeling recent observations of this phenomenon.  The closest similar efforts are those of \\cite{white01}, who present predictions from a hydrodynamic simulation at $z=1$ and $z=0.5$, and \\cite{guzik01} and \\cite{yang02}, who analyze the N-body+semi-analytic galaxy distribution of \\cite{kauffmann99a}.  One strength of the hydrodynamic simulation approach is that it predicts properties of the intergalactic medium (IGM) in addition to the galaxy population.  We have used the simulation analyzed here to model the phase distribution of baryons in the present day universe \\citep{dave01}, the X-ray background \\citep{croft01}, X-ray emission from galaxy groups \\citep{dave02}, and X-ray absorption by the diffuse IGM \\citep{chen02}. In combination with other simulations, we have used it to study cooling radiation and sub-millimeter emission from young galaxies (\\citealt{fardal01}, \\citeyear{fardal02}), the relative importance of mergers and smooth accretion in galaxy assembly \\citep{murali02}, and the correlations between galaxies and the \\lya\\ forest at high redshift \\citep{kollmeier02}.  The observed distributions and correlations of galaxy properties, in particular the luminosity function and the Tully-Fisher (\\citeyear{tully77}) relation, are also essential tests of the galaxy formation model, and any discrepancies with these observations can provide insight into the model's failings. We will consider these characteristics of the galaxy population in a separate paper (Katz et al., in preparation), though the comparison to observed galaxy-galaxy lensing measurements will bring up some of the same issues here.  We proceed to a description of our numerical methods in \\S\\ref{sec: methods}, to predictions of galaxy clustering in \\S\\ref{sec: clust}, and to the galaxy-mass correlation function and comparison to lensing measurements in \\S\\ref{sec: gdm}. We recap our findings and discuss their implications in \\S\\ref{sec: disc}. ",
        "conclusions": "\\label{sec: disc}  For the full population of galaxies with $M_b > 5\\times 10^{10}M_\\odot$, our simulation predicts a bias factor $b_\\sigma(8\\hmpc)$ of unity at $z=0$. Beneath this rms similarity of fluctuation amplitudes, however, lie numerous differences between the galaxy and dark matter distributions, in terms of present-day structure and the evolution of that structure.  1. The bias defined by rms fluctuations or by correlation functions is strongly redshift dependent.  If we choose a galaxy population of fixed comoving space density (above a redshift-dependent baryon mass threshold), then the galaxy correlation function stays approximately constant from $z=3$ to $z=0$, so that bias is initially large and declines towards unity as the dark matter catches up to the galaxies (see Figure~\\ref{fig:corrall}).  2. The bias is scale-dependent in the non-linear regime, and this scale dependence itself depends on redshift.  At $z=0$, the galaxy correlation function is depressed below the dark matter correlation function at $r\\sim 0.5\\hmpc$ and boosted above it at $r \\la 0.1\\hmpc$. As a consequence, the galaxy correlation function is well described by a power law (of slope $\\gamma = 1.78 \\pm 0.05$), while the dark matter correlation function is not (see Figure~\\ref{fig:powerlaw}). The dark matter correlation function is shallower at high redshift while the galaxy correlation function slope is roughly constant, so the rms bias factors at $2\\hmpc$ and $8\\hmpc$ start to diverge at $z\\ga 1$ (see Figure~\\ref{fig:bsigma}).  3. The correlation amplitude depends on galaxy baryon mass, with only a slight difference between the more and less massive halves of the complete sample (Figure~\\ref{fig:corrsub}b) but significant increases of $r_0$ and $b_\\sigma$ result if we select the 500 or 200 most massive galaxies in the simulation volume (Figures~\\ref{fig:r0gamma} and~\\ref{fig:bsigma}, Table~\\ref{tab:subsample}).  4. The dependence of the correlation amplitude on stellar population age is more striking than the dependence on baryon mass.  The correlation functions of the older and younger halves of the complete sample have almost identical slopes, but their correlation lengths differ by nearly a factor of two (Figure~\\ref{fig:corrsub}, Table~\\ref{tab:subsample}).  5. At $z=0$, the pairwise velocity dispersion $\\sigma_{12}(r)$ of galaxies is lower than that of the dark matter, by $\\sim 20\\%$ (Figure~\\ref{fig:pv}).  There is a large difference between the pairwise dispersions of old and young galaxies, with the preferential location of old galaxies in dense environments making $\\sigma_{12}(r)$ about twice as high at $r \\la 2\\hmpc$. The mean pairwise velocity $v_{12}(r)$ is similar for galaxies and dark matter at $z=0$, but the denser environments of older galaxies again cause a higher $v_{12}(r)$.  At $z=1$, the positive spatial bias of the galaxy population makes the mean pairwise velocity higher than that of the dark matter, though the pairwise dispersion remains somewhat lower.  6. The skewness of galaxy count fluctuations, as quantified by the hierarchical ratio $\\stf$, is lower than that of the dark matter by $\\sim 20\\%$ (Figure~\\ref{fig:s3}). The $\\stf$ ratio and its scale dependence vary with galaxy baryon mass and age, in a fairly complex way.  7. The galaxy-mass cross-correlation function depends on galaxy mass and age in much the same way that the galaxy-galaxy correlation function itself does, though the trends are weaker because the mass distribution is the same in all cross-correlations (Figure~\\ref{fig:gmxisub}).  For each class of galaxies, the bias function defined by the ratio $\\xigg(r)/\\xigm(r)$ is roughly scale-independent, even when it is significantly different from unity, though it can be somewhat different in amplitude and shape from the analogous bias function $[\\xigg(r)/\\ximm(r)]^{1/2}$ (Figure~\\ref{fig:gmxibias}).  8. The average extended mass distributions around galaxies depend significantly on baryon mass and age (Figure~\\ref{fig:massgal}). The distribution around a typical high mass galaxy falls more steeply than an isothermal distribution, so that circular velocity curves $[G\\Delta M(r)/r]^{1/2}$ decline from $r=25\\hkpc$ out to $r=1\\hmpc$.  For low mass galaxies, on the other hand, the average mass profile is shallower than isothermal at $r \\ga 100\\hkpc$, where the individual galaxy halos begin to encounter the environment of a typical surrounding group.  This effect is much more pronounced for old galaxies than for young galaxies.  The ratio of mean aperture mass $\\Delta M(r)$ to galaxy baryon mass $M_b$ is roughly independent of baryon mass for apertures ranging from $r=50\\hkpc$ to $r=1\\hmpc$, though the ratio is higher for the galaxies with the lowest $M_b$.  The ratio is higher for old galaxies than for young galaxies, again demonstrating the preferential location of old galaxies in dense environments.  For the most part, the dark matter clustering of the $\\Lambda$CDM model and these biases of the galaxy population lead to good qualitative and, where we have adequate statistics for comparison, quantitative agreement with observed galaxy clustering.  The predicted galaxy correlation function is a power law of the observed slope and approximately the observed amplitude (see further discussion below). The redshift-independence of the comoving correlation length agrees with measurements from Lyman-break galaxy surveys at $z\\approx 3$ \\citep{adelberger98,adelberger02} and from deep redshift surveys extending to $z\\sim 1$ \\citep{postman98,carlberg99}. The predicted dependence of clustering strength on baryon mass and population age agrees with the trends as a function of luminosity, color, and spectral type found in the 2dFGRS and SDSS \\citep{norberg01,norberg02b,zehavi02}. The predicted value of the hierarchical ratio $\\stf$ agrees with measurements from angular clustering catalogs \\citep{gaztanaga94,gaztanaga02,szapudi02}.  The value and scale-independence of the bias inferred from the ratio of galaxy-galaxy and galaxy-mass correlation functions agrees with recent observational estimates \\citep{hoekstra01}. The dependence of the galaxy-mass correlation function on population age parallels the dependence on galaxy type found by \\cite{mckay01}, and the aperture mass $\\Delta M(260\\hkpc)$ is approximately proportional to galaxy baryon mass, as \\cite{mckay01} also find.  There are three quantitative discrepancies between our predictions and existing measurements, two fairly subtle and one more substantial. First, the power law slope of $\\xi(r)$ for the $L_*$-galaxy sample (the 500 most massive galaxies in the box at $z=0$) is slightly too steep, $\\gamma=2.00\\pm 0.08$ compared to observed values of $1.7-1.8$, and there are hints of a departure from a power law at $r \\la 1\\hmpc$. Second, the correlation length of this sample, $r_0=4.5\\pm 0.4\\hmpc$, is low compared to the value $r_0=6.3\\pm 0.8\\hmpc$ measured for $r$-selected galaxies with $L \\approx L_*$ in the SDSS \\citep{zehavi02}, though it is consistent with the value $r_0=4.9\\pm 0.3\\hmpc$ found for $L_*$ $b_{\\rm J}$-selected galaxies in the 2dFGRS \\citep{norberg01}. Because of the still-limited size of our simulation volume, these discrepancies are of marginal statistical significance, and the N-body tests discussed in \\S\\ref{sec:corrall} imply that the low correlation length, at least, is a consequence of the lower than average fluctuation amplitude in the particular realization of $\\Lambda$CDM initial conditions used here.  The more serious quantitative discrepancy is the ratio of aperture mass $\\Delta M(260\\hkpc)$ to galaxy baryon mass $M_b$. For $M_b \\sim (0.5-1) \\times 10^{11} M_\\odot$, the simulation prediction agrees with the result derived by combining McKay et al.'s (\\citeyear{mckay01}) values of $\\Delta M/L$ with Kauffmann et al.'s (\\citeyear{kauffmann02}) stellar mass-to-light ratios, but at higher $M_b$ the predicted ratios are lower than the observationally inferred values by factors of $\\sim 1.5-2$, implying that the galaxy baryon masses in the simulation are too high (Figure~\\ref{fig:mckay}). As discussed in \\S\\ref{sec: gdm}, we believe that this discrepancy is mostly numerical in origin, since our internal resolution tests (Fardal et al., in preparation) and comparisons between different SPH implementations \\citep{springel02} suggest that galaxy baryon masses are indeed overestimated by about this factor in this mass range, relative to a simulation with higher resolution or more accurate treatment of the interface between hot and cold gas phases. A comparison between our best guess at resolution-corrected results and the \\cite{mckay01} data points yields fairly good agreement in the average values of $\\Delta M(260\\hkpc)/M_b$, which leaves little room for additional astrophysical processes, such as more aggressive feedback, to substantially suppress gas cooling and star formation in this galaxy mass range. Consistent with this conclusion, \\cite{yang02} find that the N-body + semi-analytic models of \\cite{kauffmann99a}, which do include much stronger stellar feedback, predict $\\Delta M(260\\hkpc)/L$ ratios that are too low (by a factor $\\sim 2$) relative to those of \\cite{mckay01}.  As emphasized by \\cite{guzik02} and \\cite{seljak02a,seljak02b}, the constraints from galaxy-galaxy lensing measurements complement those from the galaxy luminosity function and the Tully-Fisher and fundamental plane relations because they suffer different systematic uncertainties and have different sensitivity to astrophysical and cosmological parameters. For example, lowering $\\Omega_m$ while keeping other parameters (including the linear power spectrum) fixed would tend to lower $\\Delta M(r)/L$ while having relatively little impact on the luminosity function (though galaxy formation, and hence galaxy luminosities, would still be affected by the different timescales and baryon-to-dark matter ratios of the lower $\\Omega_m$ model). The weak lensing constraints will become considerably more powerful as the precision and detail of the measurements improves.  Where we have been able to compare, our results generally agree well with those of other hydrodynamic simulations \\citep{pearce99,pearce01,cen00,yoshikawa01}, high resolution N-body simulations that identify galaxies with sub-halos \\citep{colin99,kravtsov99}, and combined N-body + semi-analytic models \\citep{kauffmann97,governato98,kauffmann99a,kauffmann99b, benson00a,benson00b,benson01}. All of these methods make qualitatively similar predictions for the evolution of the galaxy correlation function, the bias between galaxies and mass at the present day, the dependence of that bias on galaxy type, and the relative amplitude of dark matter and galaxy pairwise dispersions. This agreement suggests that these aspects of galaxy clustering and bias arise from robust aspects of the physics of galaxy formation that all of these methods treat correctly, or at least similarly.  Our conclusions about $\\stf$ and galaxy-mass correlations are for the most part new, but we do not expect them to be fundamentally different from the predictions of other methods.  Returning to the main goals of the paper, we find that the combination of a $\\Lambda$CDM cosmology with standard ideas about galaxy formation is, on the whole, remarkably successful at reproducing observed galaxy clustering and galaxy-galaxy lensing measurements. In terms of physical interpretation, the simulation analysis offers a number of insights. It shows that the detailed relation between the galaxy and dark halo populations can account for the difference between the predicted correlation function of dark matter and the observed power law form of the galaxy correlation function, and between the predicted pairwise velocity dispersion of the dark matter and the observed, lower dispersion of galaxies. It shows that the observed dependence of galaxy clustering on galaxy properties emerges naturally from the dependence of galaxy mass and population age on environment --- basically, older and more massive galaxies form in regions that collapse early and are today biased with respect to the overall mass distribution.  It shows that values of $b_\\sigma \\sim 1$ and $\\bxix \\sim 1$ can emerge even when galaxies do not trace mass in detail.  The quantitative discrepancies discussed above may also prove instructive, as the observations improve and the uncertainties in the simulations themselves are more thoroughly understood.  Our analysis makes a number of predictions that can be tested with future data, or with future analyses of existing data. Our quantitative comparison to the observed type dependence of galaxy clustering has been limited by our finite simulation volume, which prevents us from predicting the clustering of rare classes of galaxies. However, it would be straightforward to test the quantitative predictions presented in, for example, Figure~\\ref{fig:corrsub} by taking an observed sample above a luminosity threshold that yields the same space density as our complete sample and dividing it in two based on luminosity or color (or, better still, on the basis of stellar mass and population age, using techniques like those of \\citealt{kauffmann02}). The predicted evolution of the mean pairwise velocity and pairwise dispersion (Figure~\\ref{fig:pv}) can be tested with upcoming large redshift surveys like the DEEP and VIRMOS-VLT programs. The systematic variations of $\\stf$ with galaxy luminosity and age (Figure~\\ref{fig:s3}) can be tested with the 2dFGRS and SDSS. The most informative comparisons will probably come from improved measurements of galaxy-galaxy lensing and cosmic shear, which will allow detailed tests of the trends predicted in Figures~\\ref{fig:gmxibias} and~\\ref{fig:massgal}.  On all of these fronts, observational efforts are advancing at a staggering pace.  Comparisons between more comprehensive data and improved theoretical predictions over the next few years should tell us whether our understanding of dark matter clustering and the physics of galaxy formation is indeed complete, at least in terms of features that have major quantitative impact, or whether there are still important ingredients missing from our theoretical recipe."
    },
    "0212/astro-ph0212026_arXiv.txt": {
        "abstract": "{ A catalog of optical warps of galaxies is presented. This can be considered complementary to that reported by S\\'anchez-Saavedra et al. (\\cite{sanchez-saavedra}), with 42 galaxies in the northern hemisphere, and to that by Reshetnikov \\& Combes (\\cite{reshetnikov99}), with 60 optical warps. The limits of the present catalog are: logr25 $>$ 0.60, $B_{t}$ $<$ 14.5, $\\delta$(2000) $<$ 0$^o$, -2.5 $<$ t $<$ 7. Therefore, lenticular galaxies have also been considered. This catalog lists 150 warped galaxies out of a sample of 276 edge-on galaxies and covers the whole southern hemisphere, except the Avoidance Zone. It is therefore very suitable for statistical studies of warps. It also provides a source guide for detailed particular observations. We confirm the large frequency of warped spirals: nearly all galaxies are warped. The frequency and warp angle do not present important differences for the different types of spirals. However, no lenticular warped galaxy has been found within the specified limits. This finding constitutes an important restriction for theoretical models. ",
        "introduction": "As peripheral features, disc warps are better observed in 21 cm. This has been the preferred observational technique since their discovery by Sancisi (\\cite{sancisi})  and the study by Bosma (\\cite{bosma}) until more recent samples such as those analyzed by Garc\\'{\\i}a-Ru\\'{\\i}z (\\cite{garciaruiz01a}) and Garc\\'{\\i}a-Ru\\'{\\i}z et al. (\\cite{garciaruiz01b}). Optical observations provide an important complementary tool. Even if the relation between optical and radio warps remains unclear (Garc\\'{\\i}a-Ru\\'{\\i}z \\cite{garciaruiz01a}), the great advantage of optical observations lies in the much larger samples available. Catalogs of warped galaxies have been useful to establish observational restrictions to explain warps. S\\'anchez-Saavedra et al. (\\cite{sanchez-saavedra}) first produced a catalog of 42 optical warps in the northern hemisphere out of a sample of 86 galaxies analyzed. The most noticeable result was that, taking into account the probability of non-detection of warps when the line of nodes lies in the plane of the sky, nearly all galaxies are warped, confirming the suggestion made by Bosma (\\cite{bosma}) for HI warps. Warps are therefore a universal feature, common for nearly all spiral galaxies. Even this large frequency is a severe restriction for theoretical models. As Reshetnikov and Combes said : ``Differential precession wraps any warp'', in contrast with the large frequency of real warped galaxies.  Reshetnikov \\& Combes (\\cite{reshetnikov99}) studied 540 edge-on galaxies, from which a sub-sample of 60 warped galaxies was presented. One of the most noticeable new results reported in their work was that warps were more frequent in denser environments. They also found that the warping mechanism is equally efficient in all types of spirals.  Detailed studies of particular warped galaxies in the optical, such as those by Florido et al. (\\cite{florido}) and by de Grijs (\\cite{grijs}), the latter including 44 galaxies, are also of great interest, evidently , but the production of catalogs has greater statistical power.  Here, we have examined all the galaxies contained in LEDA, with logr25 (log10 of axis ratio (major/minor axis)) $>$ 0.60, $B_{t}$ (total B-magnitude) $<$ 14.5, $\\delta$(2000) $<$ 0$^o$, -2.5 $<$ t (morphological type code) $<$ 7. The number of galaxies fulfilling these requirements is 276, which is our basis for this work, in which the galaxies were analyzed by means of the DSS. Only discs with a warp angle \\emph{wa\\/} (measured from the galactic center to the last observable point with respect to the galactic plane) larger than 4$^o$ are considered to be warped.  To determine the warp angle, \\emph{wa}, and other geometrical parameters of a warp, or even the very existence of a warp, subjective criteria were used in previous catalogs and are in turn used in the present one. Objective procedures (Jim\\'enez-Vicente \\cite{jimenez98a}, Jim\\'enez-Vicente et al. \\cite{jimenez98b}) present some problems when applied to a very large number of galaxies and were disregarded. More specifically, in some galaxies it was necessary to apply the automatic star removing procedure to an excessively large number of stars, rendering the map highly distorted. Also the separation of real warps and spiral arms was difficult to define, even for galaxies where this separation is clear to the human eye.  Our work presents 150 warps extracted from 276 edge-on galaxies. The recent study by Reshetnikov \\& Combes (\\cite{reshetnikov98}, \\cite{reshetnikov99}) extracted 60 warps from 540 edge-on galaxies. It is evident that these authors used stricter criteria to define when a galaxy is definitely warped. Our sample and scope are complementary. The Reshetnikov \\& Combes sample is limited by coordinates 0$^h$ $\\leq$ $\\alpha$ (1950) $\\leq$ 14$^h$, $\\delta$ (1950) $\\leq$ -17.5$^o$ ; ours by $\\delta$(2000) $\\leq$ 0$^o$ only. That implies that the study by Reshetnikov \\& Combes covers 17\\% of the whole sky, whereas ours covers 50\\%. This, however, is not strictly true because a large part of the Southern Sky is covered by the Zone of Avoidance due to the large extinction near the galactic plane. Taking into account this Avoidance Zone (-15$^o$ $\\leq$ b $\\leq$ 15$^o$) our sample would cover about 40\\% of the whole sky. Even if based on a smaller number of edge-on galaxies, this large coverage renders the present study more useful for certain statistical tasks, such as the distribution of the orientation of warps.  Another important difference between the work by Reshetnikov \\& Combes and the present study is that the former only pays attention to the discs of spiral galaxies while ours includes lenticular galaxies. This inclusion is very important from the theoretical point of view because, roughly speaking, lenticulars have a disc but not gas; in other words, the distribution of stars in lenticulars is very similar to that in spiral galaxies. This is an overgeneralisation, because even if the amount of gas in lenticulars is less than in even late-type spirals, lenticulars do process some gas in the outer parts. Another important difference between lenticulars and spirals is that the former have a dominant bulge. This fact makes analysis less simple, as huge bulges could in principle affect the formation of warps. In any case observations of warps in lenticulars have yet to be made and may introduce decisive restrictions on the mechanisms behind warps.  A large number of theories can be found in the literature: non-spherical dark halos misaligned with the disc (Sparke \\& Casertano \\cite{sparke}; Dubinski \\& Kuijken \\cite{dubinski}), gas infall into the dark matter halo (Ostriker \\& Binney \\cite{ostriker}; Binney \\cite{binney}) or directly into the disc (L\\'opez-Corredoira et al. \\cite{lopez-corredoira}), interactions with companions or small satellites (Weinberg \\cite{weinberg}), intergalactic magnetic fields (Battaner et al. \\cite{battaner}), etc. Kuijken \\& Garc\\'{\\i}a-Ru\\'{\\i}z (\\cite{kuijken}) recently presented a concise review summarizing several mechanisms proposed to explain warps.  The large fraction of warped galaxies seems to exclude one of the most obvious models, based on interactions, but this hypothesis has been reconsidered by Weinberg (\\cite{weinberg}), assuming a strong tidal amplification by the gravitational wake of a satellite, although this fact was not confirmed by Garc\\'{\\i}a-Ru\\'{\\i}z (\\cite{garciaruiz01a}) who also argued that the warp of the Milky Way induced by the Magellanic Clouds should have been predicted to have a very different orientation from that observed. Warps are common in merging systems (Schwarzkopf \\& Dettmar \\cite{schwarzkopf}) but it remains unclear whether mergers or interactions can explain all, or at least most, warps.  Warps are more frequent in denser environments (Reshetnikov \\& Combes \\cite{reshetnikov99}). Early $z \\approx 1$ warps were considered by Reshetnikov et al. (\\cite{reshetnikov2002}). Early warps were larger, which favors the interaction model.  Other models cannot be completely excluded from the observation of early $z \\approx 1$ warps. Magnetic fields were also much stronger and the rate of infall of matter onto a galaxy was higher.  Garc\\'{\\i}a-Ru\\'{\\i}z (\\cite{garciaruiz01a}) observed that, even if galaxies with extended discs may be warped, extended discs are less frequent in denser environments.  The observational study of warps in lenticular galaxies is crucial. If warps are absent or are less common in lenticular galaxies which are gas-poor, those models based on gravitation alone would have difficulty in explaining this fact. Models in which a permanent torque acts on the gas of the galaxy would have the preference. For instance, models like that by Kahn \\& Woltjer (\\cite{kant}) or its more recent version by L\\'opez-Corredoira et al. (\\cite{lopez-corredoira}) would be favored. Note that neither model requires the assumption that galaxies have a large dark matter halo. The magnetic model, in which intergalactic magnetic fields maintain the warp structure, would acquire additional support from this fact.  It therefore seems that the compilation of large samples of warped galaxies, even if they contain just a few parameters about their position, the parent galaxy and the magnitude and shape of the warping, contributies as much as the detailed examination of the HI maps of a few galaxies. ",
        "conclusions": "\\begin{enumerate} \\item This catalog contains a large sample of warped galaxies (150), covering the complete southern hemisphere. We present the whole sample of 250 spirals and 26 lenticulars, with limits logr25 $>$ 0.60, $B_{t} <$ 14.5, $\\delta$ $<$ 0$^o$, -2.5 $<$ t $<$ 7. It is especially suitable for statistical analysis and, indeed, has already been used to study the relation between intrinsic parameters and warps by Castro-Rodr\\'{\\i}guez et al. (\\cite{nieves}). The catalog may obviously be used to choose which galaxies are to be observed in detail. \\item There is the unavoidable problem of the existence of other contaminant effects that could present the appearance of warps. Reshetnikov \\& Combes (\\cite{reshetnikov98}) estimated that about 20$\\%$ of the features assumed to be warps could actually be spiral arms. This value could be applied here. This effect introduces small errors into our calculated frequencies, which should be taken into account in statistical studies. \\item We confirm that warps appear to be a universal feature in spiral galaxies. The frequency of warps is very high (60\\%), and because of the difficulty (or impossibility) of detecting warps with the line of nodes in the plane of sky, it is concluded that most (if not all) galaxies are warped. Bosma (\\cite{bosma}), S\\'anchez-Saavedra et al. (\\cite{sanchez-saavedra}) and Reshetnikov \\& Combes (\\cite{reshetnikov99}) previously reached this conclusion. \\item We also confirm, by means of a larger amount of data, the finding by Reshetnikov \\& Combes (\\cite{reshetnikov99}) that warps are equally present in all types of spirals. The distribution of warps for the different types of spirals coincides with the distribution of galaxies with type. The maximum observed warp angle either from the center (\\emph{wa}) or from the starting radius of the warps ($\\beta$) has no relation with the type of spiral. We have observed that very large values of $\\beta$ do not correspond to large warp angles of \\emph{wa}. Shorter warps are steeper. \\item We found no warped lenticular galaxies at all among the 26 galaxies within our limits, with logr25 $>$ 0.60. We enlarged the sample to reach logr25 $>$ 0.57 and again, none of the 38 lenticulars was warped. There is a sudden transition at t=0. All spirals are warped; no lenticular is warped. The main difference between spirals and lenticulars is probably that the former are gas rich and the latter gas poor. Gas seems to be a necessary ingredient in the warp mechanism. Models based on gravitation alone would have serious difficulties in explaining this. \\end{enumerate}"
    },
    "0212/astro-ph0212210_arXiv.txt": {
        "abstract": "Interferometers require accurate determination of the array configuration in order to produce reliable observations.  A method is presented for finding the maximum-likelihood estimate of the telescope geometry, and of other instrumental parameters, astrometrically from the visibility timelines obtained from observations of celestial calibrator sources. The method copes systematically with complicated and unconventional antenna and array geometries, with electronic bandpasses that are different for each antenna radiometer, and with low signal-to-noise ratios for the calibrators. The technique automatically focusses on the geometry errors that are most significant for astronomical observation. We apply this method to observations made with the Very Small Array and constrain some $450$ telescope parameters, such as the antenna positions, effective observing frequencies and correlator amplitudes and phase shifts; this requires only $\\sim 1$ h of CPU time on a typical workstation. ",
        "introduction": "Accurate knowledge of observing frequency, projected baselines and other geometric parameters is essential if interferometer data are to be turned into useful images. The accuracy required is typically a small fraction of the operating wavelength. For example, to keep phase errors below 10 degrees for an interferometer operating at 30~GHz, it is necessary to know the baseline length to an accuracy of better than 0.3~mm.  The standard way in which interferometer geometries are determined is along the following lines. First, direct measurements are made using, for example, tape measures, theodolites, and measurements of electrical lengths of electronic systems. Then, observations of celestial calibrators are carried out. For example, the observation by a particular pair of antennas of a bright calibrator of known RA and Dec gives the collimation phase of that pair (the difference in electrical length of the two arms of that interferometer).  How this quantity varies, for example, with hour angle and for a series of calibrators at different declinations, is also investigated. These standard methods are described in e.g. \\citet{thompson94}, and in these ways a series of corrections to the directly measured geometry is produced, with the aim of ensuring that the physical baselines (in wavelengths) used in subsequent data analysis are correct.  This approach is of proven effectiveness for determining the geometries of interferometer arrays such as the Very Large Array (VLA). It is not adequate, however, for interferometers such as the Very Small Array (VSA, see e.g. \\citet{watson02}; \\citet{scott02}), which is designed for imaging the CMB on scales $\\ga 10$ arcmin at 30~GHz, and thus has baselines of length $\\la 2$~m. At first sight this may seem surprising, as VSA baselines (in both metres and wavelengths) are so small. From the point of view of geometry determination from astrometry, however, the VSA is different from most other interferometer arrays in two respects. \\begin{itemize} \\item The collecting area of VSA antennas is very small and so there are few suitable celestial calibrators. Only the brightest calibrators are of use and signal-to-noise is a major issue. \\item For most arrays like the VLA, the principal swivel axes of each antenna (e.g. altitude and azimuth, or polar and declination) intersect, at least nominally. If an axis is incorrectly aligned (e.g. the polar axis does not quite point to the Pole), then to first order this has no effect on the measured phase of a source, and its effect on source amplitude will manifest as a simple pointing correction which a standard 5-point observation (at expected source position, then offsets North, South, East then West) will provide. In the case of the VSA, the table elevation axis and the individual horn rotation axes (see Section~\\ref{sec:geovsa}) do not intersect by $\\sim$ 1~m. This design helps reduce important systematic errors and also reduces cost, but it makes astrometric geometry determination significantly harder. If an axis is misaligned, then in general a phase error is also introduced. The identification of such errors is exacerbated by the lack of signal-to-noise. \\end{itemize} There are also further difficulties associated with determining the geometry of the VSA, which can be present in other interferometers. \\begin{itemize} \\item As a result of, for example, the compliance needed for fitting horn-reflectors to cryostats, the position of the electrical centre of each horn will be slightly different, manifesting as a physical baseline correction and a pointing correction. \\item A large observing bandwidth makes it hard to achieve a flat band pass, so the effective radio frequency becomes antenna dependent. \\item Unconventional geometry makes it difficult to make a straightforward interpretation of a set of geometry calibration observations. \\end{itemize}  As a result of all these factors, phase (and amplitude) effects arise from complicated causes that are particularly hard to determine because of the lack of very high signal-to-noise calibrators. We have therefore implemented a maximum-likelihood method to determine the geometry, and other telescope parameters, astrometrically from the visibility timeline data obtained during a calibration observation of a sufficiently bright unresolved radio source. In the geometric model, we include as variables all the important dimensions and angles that require accurate determination, including the positions of the antennas, their effective observing frequencies and phase shifts on different baselines. In finding the best-fit geometry, the maximum-likelihood method correctly allows for the signal-to-noise ratio of the observations (on the assumption that the instrumental noise is Gaussian); the procedure also automatically guards against attempting to determine unimportant geometry errors that do not in fact much affect astronomical observation, since these similarly will not much affect calibrator observations.  In this paper, we focus on the method in its application to the VSA, although we point out the usefulness of the approach generally to radio and optical interferometry. In Section~\\ref{sec:geocal}, we present the general maximum-likelihood method for the geometry calibration. The VSA is introduced in Section~\\ref{sec:geovsa}, where we develop a model of the telescope as a function of the parameters to be calibrated.  Using this model, the likelihood method is applied to real VSA observations in Section~\\ref{sec:geoapplications}.  Finally, our conclusions are presented in Section~\\ref{sec:geoconclusions}. ",
        "conclusions": "\\label{sec:geoconclusions}  In this paper, we have presented a method for astrometric geometry calibration of interferometers by performing a likelihood analysis of visibility timelines produced during an observation of a reasonably bright radio point source. The method has been demonstrated for the specific case of the Very Small Array, for which one must simultaneously constrain some 450 telescope parameters, such as the antenna positions, the effective observing frequency of each antenna which can vary due to different bandpass characteristics, and phase shifts on different baselines introduced in the electronic system or by different cable lengths. In particular, analysis of VSA observations of Tau-A and Cas-A sources show that the relative positions of the antennas on the telescope table can be determined to a small fraction of a wavelength. The quality of the fit is sufficiently good to enable high dynamic-range mapping with the VSA. This astrometric calibration is very straightforward and cheap, as it requires no equipment other than the telescope itself.  Multiple calibration observations can easily be combined to produce a better parameter estimate by summing the $\\chi^2$-functionals for all available data sets.  The calibration method is currently in active use by the VSA project and has successfully determined the geometry for several different antenna configuration since the start of its observing programme.  The calibration of the VSA is described in more detail in \\citet{watson02}, and power spectra derived from the first year of observations in the compact configuration are presented in \\citet{scott02}.  Although the astrometric calibration method has been applied specifically to the VSA, it can easily be adapted to other telescopes. It is worth noting, however, that there is one feature of the VSA which does not necessarily generalise to other CMB telescopes, even though it is common to most ground-based arrays.  Each VSA antenna tracks the source separately in `azimuth' ({\\em semi-comounted} design).  This has the advantage of providing a valuable check on systematic contaminations, because the path delays on a particular baseline change during the observation and the sky signal can be recognised by its characteristic fringe frequency.  On the Cosmic Background Imager \\citep{padin02} and the Degree Angular Scale Interferometer \\citep{leitch02}, the antennas are mounted on a rigid tracking platform supported by an alt-azimuth mount.  This design keeps the baseline projections with respect to the source constant and reduces the fringe rate on all baselines to zero, making the described geometry fitting impossible.  A similar method can still be applied, however, if the fringe rate is artificially changed.  For instance, the telescope could be kept at a fixed pointing position while the source is within the antenna beam, thus providing a non-zero fringe rate."
    },
    "0212/astro-ph0212576_arXiv.txt": {
        "abstract": "In the first three years of operation STIS obtained slitless spectra of $\\sim 2500$\\ fields in parallel to prime HST observations as part of the STIS Parallel Survey (SPS).  The archive contains $\\sim 300$\\ fields at high galactic latitude ($|b|>30$) with  spectroscopic exposure times greater than 3000 seconds. This sample contains 220 fields (excluding special regions and requiring a consistent grating angle) observed between 6 June 1997 and 21 September 2000, with a total survey area of $\\sim 160$\\ square arcminutes.  At this depth, the SPS detects an average of one emission line galaxy per three fields.  We present the analysis of these data, and the identification of 131 low to intermediate redshift galaxies detected by optical emission lines.  The sample contains 78 objects with emission lines that we infer to be redshifted [OII]3727 emission at $0.43<z<1.7$. The comoving number density of these objects is comparable to that of \\ha-emitting galaxies in the NICMOS parallel observations. One quasar and three probable Seyfert galaxies are detected.  Many of the emission-line objects show morphologies suggestive of mergers or interactions. The reduced data are available upon request from the authors. ",
        "introduction": "The relative pointing of instruments on the Hubble Space Telescope allows for parallel observation during targeted science exposures.  Parallel fields are essentially random pointings on the sky, being separated by 5-8 arcminutes from the prime target. Extensive imaging of parallel fields was possible with WFPC2, and the Medium Deep Survey obtained hundreds of fields (c.f. Griffiths et al.\\ 1994). With the addition of the slitless spectroscopic capabilities of STIS (Kimble et al.\\ 1998) and NICMOS (Thompson et al.\\ 1998), parallel observations now allow HST to conduct redshift surveys in large numbers of random fields without negatively impacting the duty cycle of the telescope (see Gardner et al.\\ 1998, McCarthy et al.\\ 1999, Yan et al.\\ 1999).  These untargeted redshift surveys explore a new regime in galaxy evolution, as they are less subject to a bias induced by cosmic variance than pencil beam surveys (e.g. Cohen et al. 2000) and detect fainter objects than wide field surveys (e.g. CFRS; Hammer et al.\\ 1997).  The slitless spectroscopy requires comparable deep imaging for object identification; thus the redshift survey is accompanied by high resolution imaging.  In this paper, we present the analysis of the deepest spectroscopic fields in the STIS parallel survey (SPS).  STIS parallel spectra were taken over the 5300-10000 \\AA\\ range during the course of the program.  We have searched for emission line galaxies, in particular [OII]-emitting galaxies at redshift $0.43<z<1.7$.  Our motivation is to determine the comoving density of star-formation at these redshifts as inferred from [OII] emission (Teplitz et al.\\ 2002; hereafter Paper II). We will discuss the target selection in section 2, the data reduction, line identification and measurement in section 3, and the results in sections 4 and 5. ",
        "conclusions": "The deep SPS fields cover 160 square arcminutes (with 141 square arcminutes of usable area) along 219 distinct lines of sight.  131 objects with emission lines were identified.  Forty-four fields have multiple emission line objects.  Emission-line redshifts are inferred ranging from z=0.034 to z=1.548.  Seventy-eight objects are identified as having [OII] emission, of which 13 also have other emission lines. Forty objects are identified primarily with \\ha\\ emission, and 13 with [OIII] or \\hb\\ in the absence of \\ha\\ and [OII].  Of these, 37 have secure redshifts (multiple lines; quality 4), 25 have quality 3, 28 have poor redshifts (quality 1), 15 have problem redshifts (quality 1.5 or 2), and 26 appear to have \\ha\\ emission for which we do not measure line fluxes.  Figure \\ref{fig: zhist}\\ shows the distribution of inferred redshifts. The detection of objects falls off at the edges of the wavelength range as the sensitivity drops.  At the lowest redshifts, the falloff in \\ha\\ objects is attributable to the small volume covered by the survey.  Figure \\ref{fig: maghist}\\ shows the distribution of apparent magnitude.  A rough calculation of the density of detected objects can be made by considering the [OII]-emitters detected at the redshifts that place the [OII] line in the good regions of the sensitivity curve, ($0.5<z<1.2$).  The SPS surveys a comoving volume of $1.62\\times 10^5$\\ $h_{50}^{-3}$\\ Mpc$^3$\\ for $\\Omega_M=1, \\Omega_{\\Lambda}=0$ (or $1.56\\times 10^5$\\ $h_{70}^{-3}$\\ Mpc$^3$\\ for $\\Omega_M=0.3, \\Omega_{\\Lambda}=0.7$) between those redshifts, and detects 63 [OII]-emitters, giving a comoving number density of $3.7\\times 10^{-4}$ $h_{50}^{-3}$\\ Mpc$^{-3}$.  This density is slightly higher than that detected by the NICMOS parallels at $0.7<z<1.9$\\ (McCarthy et al.\\ 1999). As McCarthy et al.\\ (1999) point out, this density is lower than the density of galaxies brighter than $L^*$\\ today (Ellis et al.\\ 1996), not all of which are star-forming, but is comparable to that of Lyman Break Galaxies at $z\\sim 3$\\ (Steidel et al.\\ 1996). The SPS sample of [OII] emitters is sufficient to measure the comoving star-formation density at a median $z\\sim 0.9$, and we will present that analysis in Paper II.  The SPS fields provide promising results for the detection of intermediate redshift galaxies in optical slitless spectroscopy of random fields.  While the survey is subject to a number of biases (towards compact objects and away from low equivalent width lines), it is an efficient means to detect distant, faint galaxies.  The high resolution of the direct images should also enable a study of the morphological properties of intermediate redshift galaxies.  The installation of the Advanced Camera for Surveys (ACS; Ford et al.\\ 2001) on HST will make a more extensive survey possible in the coming years.  With its larger field of view and filtered imaging capabilities, ACS will refine the procedures developed for the SPS program."
    },
    "0212/astro-ph0212095_arXiv.txt": {
        "abstract": "{The latitudinal sound-speed structure of the Sun's convection zone gives insight into the physical processes occurring there, specifically the cellular convection and possibly the presence of magnetic fields. Using helioseismic data from the GONG network and MDI instrument on SOHO, we map the latitudinal acoustic structure of the convection zone from 1995 to 2002. The temporally averaged structure confirms previous findings of an excess in sound speed at the $10^{-4}$ level at 60 degrees latitude. There also appear to be some variations with time, with the peak in sound-speed asphericity at 60 degrees growing towards the maximum of solar activity according to the MDI data. However, we present some evidence that such variation may be associated with instrumental variation between the epochs before and after SOHO was temporarily lost in 1998. Nonetheless, some genuine temporal variation may be present, and we discuss the possible physical causes of that. ",
        "introduction": "Analysis of the global p-mode oscillations of the Sun reveals the Sun's internal structure and constrains the physics of the solar interior. In this paper we focus on the latitudinal variation of the structure of the solar convection zone and the possible temporal variation of that structure during the rising phase of solar cycle 23. The convection zone occupies the outer thirty per cent by radius of the Sun, as was deduced from helioseismic inversions for the Sun's radial hydrostatic structure (Christensen-Dalsgaard, Gough \\& Thompson 1991; cf., Gough 1977; Ulrich \\& Rhodes 1977). Such an investigation uses the mean multiplet p-mode frequencies to deduce the spherically symmetric component of the Sun's structure. But with high-quality data such as from the Global Oscillation Network Group (GONG) and the Michelson Doppler Imager (MDI) instrument on board the SOHO satellite (Harvey et al.~1996; Scherrer et al.~1995) it has become possible to make deductions also about the latitudinal variation of the structure from the dependence on the azimuthal order of the frequencies within multiplets (Gough et al.~1996; Antia et al.~2001). What is measured is acoustic asphericity, i.e., the deviation from spherical symmetry of the propagation speed of the p-mode waves as a function of position. The relevant wave speed is the adiabatic sound speed but possibly modified by magnetic effects in regions of magnetic fields. The acoustic structure of the convection zone may be expected to reflect the asphericity caused by the Sun's rotation, but also perhaps the cellular convection and the presence of magnetic fields. At the resolution scale at which the global helioseismic inversions can be undertaken, we are not here concerned with granular or even supergranular scales. But the thermal effects of giant cells or other such large-scale circulatory structures may be detectable in the inferred latitudinal variation of sound speed, as perhaps may also be a sufficiently strong large-scale magnetic field. In the present paper we present results of inversions of a more extensive data set than was considered by Antia et al.~(2001). We also show that there is some evidence for the temporal variation of the latitudinal structure during the rising phase of the solar activity cycle. However, one should be cautious because of possible instrumental variation during the period of observations, and so we wish to examine such a possibility more closely.  \\begin{figure*}[t] \\hbox to \\hsize{\\resizebox{\\figwidth}{!}{\\includegraphics{ms3124f1a.eps}} \\quad\\hfill \\resizebox{\\figwidth}{!}{\\includegraphics{ms3124f1b.eps}}} \\caption{Contour diagram of the temporal average of aspherical component of squared sound speed as a function of radial distance and latitude using the GONG (left-panel) and MDI (right-panel) data. The MDI data used are restricted to $\\ell < 110$ (see Section 4). Solid and dotted contours denote positive and negative values respectively. Contours are drawn at interval of $2\\times10^{-5}$.} \\label{fcs} \\end{figure*} ",
        "conclusions": "The helioseismic inversions for asphericity in sound speed show a distinct peak at a latitude of $60^\\circ$ inside the convection zone. These inferences raise a number of questions. What is so special about $60^\\circ$ latitude? Perhaps it is the boundary between two oppositely directed meridionally circulating cells, or at least the poleward limit of a low-latitude cell. If the perturbation to the sound speed is indeed caused by magnetic field, as we have suggested (Antia et al.~2000), then the flux could accumulate at this latitude in a convergent lane between the two cells. We can envisage two possibilities. Earlier studies of surface magnetism emphasised the poleward migration of prominences and faculae from around this location (e.g., Waldmeier 1981; Makarov \\& Sivaraman 1989a,b) and led to a picture where this latitude had surface flows diverging from it. In that case, the cells could act to dredge dynamo-generated flux from the lower convection zone into this latitude. On the other hand, ring diagram analyses of local helioseismic data (Basu et al.~1999; Haber et al.~2000) find poleward meridional circulations at low latitudes, in the outer 20 Mm or so that they can image. Small-scale fields could be generated by the alpha effect operating in the upper convection zone, or could come from the decay of old sunspots and active regions. The meridional flows could sweep up this near-surface flux and collect it like jetsam in a downflow lane at $60^\\circ$ latitude.   A further point of interest about $60^\\circ$ latitude is that this is where the near-surface shear changes sign, according to MDI helioseismic measurements: the angular velocity increases with depth beneath the surface at low latitudes whereas it decreases with depth at high latitudes (Schou et al. 1998; Corbard \\& Thompson 2002). This could reflect the role of differently signed meridional circulation flows in the transport of angular momentum on either side of $60^\\circ$. Furthermore, the time-dependence of rotation at high latitudes is perhaps another manifestation of the interplay between magnetic field and flows. The variation in rotation rate in the upper convection zone shows a pattern with bands of faster and slower than average rotation rate moving towards the equator with time at low latitudes (Howe et al.~2000; Antia \\& Basu 2000) while at high latitudes these bands appear to be moving towards the pole (Antia \\& Basu 2001). The transition between equatorward and poleward movement occurs around a latitude of $50^\\circ$. The near-surface rotation at high latitudes, as revealed by helioseismology, has been slowing down until 1999 during the rising phase of the present cycle (Schou 1999; Antia \\& Basu 2001). At high latitudes, the rotation rate decreases with time, reaches a minimum and then starts increasing again. The epoch at which the minimum occurs is a function of latitude: it gradually shifts to later times as one goes to higher latitudes. One might envisage that it is the open magnetic field lines that slow the rotation. As the open flux retreats glacier-like back to high latitudes, it releases its brake on the rotation at a given latitude and so at that latitude the rotation rate springs back up.    One of the outstanding problems in solar physics is to identify the location of a plausible mechanism responsible for simultaneous variations of the Sun's luminosity, radius, effective temperature and p-mode frequencies with the activity cycle. The seat of the solar dynamo is generally believed to be located in the tachocline region beneath the convection zone where the mechanism for solar luminosity variation could also reside. However, it has been argued by Balmforth et al.~(1996) that an agent like a magnetic field concentrated at the base of the solar convection zone is unlikely to be directly responsible for the observed cyclic modulations; nor can the purely thermal changes occurring in the near-surface layers of the Sun account for the observed irradiance variations. Any deeply seated driving mechanism would yield values of the quantity $W = (dR/R)/(dL/L)$, which measures the relative changes in solar radius and luminosity, somewhere between $0.2$ and $0.5$. The current estimates of $W$ seem to suggest much lower values, of order $0.02$ (Emilio et al.~2000). In fact, these results from satellite observations indicate that the principal cause of the solar luminosity variation is unlikely to be located very deep inside the Sun.  Helioseismic inversions have revealed the existence of a shear layer below the photosphere, extending to depths of about 50 Mm below the photosphere, with its peak occurring at 30 Mm depth, in the midst of the helium ionisation zone. (The above values depend somewhat on latitude, the depth of shear layer reducing with increasing latitude.) Dikpati et al.~(2002) have explored the role of this subphotospheric radial shear layer in driving a solar dynamo, and find that a differentially rotating shear layer is capable of generating toroidal magnetic fields of up to 1 kG strength. Antia et al.~(2000) had invoked the pumping of downward moving plumes (downdrafts) to concentrate toroidal magnetic field in the outer convection zone and to enhance its strength to values approaching 70 kG at depths around 70 Mm below the surface. This takes place at depths around 70 Mm below the surface, where some sort of equipartition is achieved between the magnetic energy density $\\langle B^2 / 8\\pi\\rangle$ and the kinetic energy density of downdrafts $\\langle \\rho w^2\\rangle$. Note, the velocities of the order of 1 km/s adopted here, refer to the downdraft motions which have velocities several times those given by the standard mixing length theory for convective elements (cf., Mestel 1999). This process, of course, does not exclude the generation of a large fraction of the toroidal magnetic field in the solar interior, with the help of differential rotation in the tachocline. We are merely pointing out that part of the toroidal field in the convection zone may be generated in the sub-surface shear layer itself, where the poloidal field could be reinforced by some kind of alpha effect. This would then be immediately available for further generation of the toroidal field by the action of strong subsurface radial shear.  Earlier it was noted by Antia et al.~(2000) and Dziembowski et al.~(2000) that residual splitting coefficient $a_2$ showed a striking peak around $r=0.96R_\\odot$. If this peak should result solely from the presence of large-scale magnetic field at depths of order 30 Mm, the inferred magnetic field strength comes out to be 20--30 kG. But with the removal of surface term, this peak shifts downwards to 70 Mm depth and required magnetic field strength would be about 70 kG.  It is readily seen from Fig.~1 that the contours of constant $\\delta c^2/c^2$ in both the GONG and MDI data show a pronounced radial extension around the $60^\\circ$ latitude peak situated at a depth of about 70 Mm. Even though a comparison of the contours of constant $\\delta c^2/c^2$ in the pre-1998.4 and post 1999.4 GONG and MDI data sets displayed in Fig.~8 barely shows a detectable temporal variation, the difference in acoustic asphericity between the high and low activity averages shown in the bottom panels of Fig.~8 has some indication of a possible temporal variation of asphericity around a depth of $0.1R_\\odot$ at both equatorial and higher latitudes. It may therefore be tempting to speculate that there is a reservoir of magnetic energy stored in the outer $0.1R_\\odot$ of the convection zone which perhaps undergoes some temporal variation. Such a time-varying storage of magnetic energy might even play a role in driving the solar activity cycle. We suggest that the solar irradiance variation is probably due to the periodic release of such a stored magnetic energy in the outer layers of the Sun, resulting from an interplay between the solar magnetic field, subsurface radial shear and pumping by convective downdrafts."
    },
    "0212/astro-ph0212430_arXiv.txt": {
        "abstract": "The radio source 3C~84, in NGC~1275, has a two sided structure on parsec scales.  The northern feature, presumed to be associated with a jet moving away from the Earth, shows strong evidence for free-free absorption.  The ionized gas responsible for that absorption would be a source of detectable stimulated recombination line emission for a wide range of physical conditions.  The VLBA has been used to search for the H65$\\alpha$ hydrogen recombination line.  The line is only expected to be seen against the northern feature which contains a small fraction of the total radio flux density.  This spatial discrimination significantly aids the search for a weak line.  No line was seen, with upper limits of roughly 15\\% of the continuum over a velocity range of 1486 \\kms\\ with resolutions up to 6.6 \\kms.  In the absence of a strong radiation field, this would imply that the free-free absorbing gas has a wide velocity width, is moving rapidly relative to the systemic velocity, or is concentrated in a thin, high density structure.  All of these possibilities are reasonably likely close to an AGN.  However, in the intense radiation environment of the AGN, even considering only the radiation we actually observe passing through the free-free absorbing gas, the non-detection is probably assured by a combination of saturation and radiation damping. ",
        "introduction": "The radio source 3C~84 is associated with NGC~1275, the dominent member of the Perseus Cluster.  The radio source has a complex structure on parsec scales that has been an object of study with VLBI since the very early days of that technique (see Walker et al. 2000 and references therein).  There is a bright core region, compact at high frequencies, which is thought to be close to the central engine of the AGN.  Extending to the south, currently to about 15 milliarcseconds (mas), is a complex structure clearly related to the jet seen on larger scales, but showing a morphology reminiscent of a radio lobe.  That structure has been moving outward at approximately 0.3 \\masr ~and its size extrapolates to zero at about 1959 when 3C~84 started to increase in flux density from below 10 Jy to values in excess of 50 Jy reached in the 1970's and 1980's.  The flux density has been decreasing lately, but 3C~84 is still one of the brightest compact sources in the sky.  HI absorption is seen against the nuclear region at both the optical systemic velocity of $v=5260$ \\kms\\ \\citep{dV91} and at the velocity of a possibly infalling system at about $v=8200$ \\kms.  VLBI observations of this absorption have provided information on the intervening clouds \\citep{M02} which are thought to be located along the line-of-sight far from the nucleus.  In 1993, a structure extending about 8 mas north of the core was found in early Very Long Baseline Array \\citep[VLBA;][]{N94} observations at 8.4 GHz \\citep{W94} and in Global VLBI Network observations from 1991 at 22 GHz \\citep{V94}.  By assuming that the northern and southern features originated in the core at the same time, a simple relativistic beaming model gives an intrinsic jet velocity of about $0.4c$ and an angle to the line of sight of about $40\\cdeg$ \\citep[][updated for $H_o=75$ \\kms Mpc$^{-1}$, the value that is used in this paper]{W94}.  The assumption of a common start time is reasonable given the flux history of the source and the measured velocities.  The northern feature was much brighter at the higher frequency, suggesting that the radiation is free-free absorbed. The absorption is only seen against the counterjet.  The near-side jet is not absorbed at all and the core is at most weakly absorbed; synchrotron self-absorption may be responsible for the inverted spectrum.  Therefore, the absorbing material is most likely located on parsec scales ($1\\ {\\rm pc} = 3\\ {\\rm mas}$), probably associated with the presumed accretion disk that feeds a central black hole.  The implications of the observed absorption for disk structure are discussed by \\citet{L95}.  The counterjet absorption was studied in detail with near-simultaneous VLBA observations at many frequencies in 1995 by \\citet{W00}.  It was also seen on somewhat larger scales at lower frequencies by \\citet{S98}.  These observations left little room to doubt that free-free absorption is the explanation of the observed spectrum of the counter jet.  They also were able to determine the two-dimensional distribution of the absorption over the region of the counterjet. That structure was found to be dominated by a radial gradient decreasing with core distance.  If the gradient is fitted with a power law, the index is a bit above $-2$.  But a power law is not an especially good description of the observed gradient.  If a temperature of $10^4$ K is assumed, the absorption implies a gas with an emission measure (EM) of about $5\\times 10^8$ \\emu\\ at a projected distance of 2.5 pc from the core.  The theory of radio recombination lines predicts that weak maser-like emission can occur over almost the entire radio frequency range due to non-LTE effects \\citep[see, for example,][]{S76}.  The strength of the lines, determined by the degree of inversion of the population levels, is a sensitive function of density and temperature.  Calculations of the expected strength of recombination lines for gas parameters implied by the free-free absorption and geometry in 3C~84, in a benign radiation environment, showed that there was a significant chance of being able to detect such lines.  An observation of such a line would provide significant constraints on the physical conditions in the absorbing gas.  The calculations indicated that the strongest lines would be those in the general vicinity of 20 GHz.  The observation of recombination lines from the counterjet region of 3C~84 is tricky.  The counterjet only provides a small fraction of the total flux density, so, even if the optical depth against that region is moderately high, the effect on the total power spectrum would be small.  This would place severe requirements on the bandpass calibration of observations that don't resolve the source.  We have used the VLBA, along with one antenna of the VLA, to attempt to detect the H65$\\alpha$ line ($\\nu_{rest}=23.404284$ GHz) in 3C~84.  With the resolution of the VLBA, the region where the line is expected to be seen can be separated spatially from the rest of the source, relaxing the bandpass calibration requirements.  And the fact that the recombination line is not expected to be seen against the bright parts of the source allows those regions to be used to do an accurate bandpass calibration using a method developed for these observations.  The calculations of expected line strengths that encouraged us to make the observations reported here ignored the possible effects of the intense radio radiation from the jets.  It turns out that those effects can be severe.  \\citet{S78} showed that a strong background radio source can alter the level populations and saturate the recombination line maser.  With saturation, the amount of amplification is limited by the pump rate, so the line strength can be much lower than would be calculated without taking saturation into account.  The conditions in 3C~84 are such that the system is likely to be at or near saturation.  Even more serious, \\citet{S78} pointed out that radio recombination lines will be severely broadened by radiation damping in the presence of a strong radio source.  Inserting numbers appropriate for 3C~84 into equation 24 of \\citet{S78} gives a line width of thousands of \\kms.  This would make the line unobservable.  Therefore detection of a line would imply either that the absorbing region is at least twice as far from the radio source than implied by the geometry described above, or that something is wrong with the physics.  Likewise, a non-detection, unfortunately, provides no interesting constraints on the physical conditions in the 3C~84 accretion region. ",
        "conclusions": "As noted in the Introduction, the radio continuum observations give good evidence that the parsec scale structure of 3C~84 includes an opposing pair of jets inclined at about $40\\cdeg$ to the line-of-sight with significant amounts of ionized gas in the region between the jets.  The ionized gas is most likely associated with the accretion disk that would be there in most AGN models.  The projected distance from the northern feature to the core is about 2.5 pc so, with the above geometry, the region of the jet that emits the absorbed radiation is about 5 pc from where the radiation passes through the disk on the way to the Earth.  That is a small enough distance that, as also noted in the Introduction, gas at the disk would be subject to a sufficiently high radiation intensity that radiation damping would be expected to broaden the H65$\\alpha$ recombination line to thousands of \\kms, preventing detection in our observations. Our non-detection is consistent with this expectation.  If a line had been detected, the lack of excessive radiation damping would constrain the gas to lie beyond some minimum distance from the radio emitting regions.  The actual distance limit depends on which portions of the overall radio source illuminate the absorbing gas. The bare minimum illumination is from the northern feature, the feature for which the free-free absorption is seen.  If that is the only radiation illuminating the absorbing gas, the separation between the northern jet and the absorbing region would have to be at least 11 pc for the radiation damping line width to be below the approximatly 500 \\kms\\ upper limit for what could have been detected in our observations.  That is over twice the distance between the northern jet and the disk in the geometry derived from the continuum observations.  It is a bare minimum distance, because both the southern feature (near side jet) and the core region most likely also illuminate the absorption region.  With this additional illumination, the minimum distance between the emitting and absorbing regions would need to be much higher, up to about 50 pc if all regions illuminate the absorbing gas.  Thus if a line had been observed, and the physics is right, the absorbing region would have to be sufficiently far from the various radio sources that the idea that it is associated with the accretion disk would be wrong.  That, in turn, would leave us struggling to explain why only the northern feature is free-free absorbed.  \\placefigure{recomb10000}  The free-free absorption of the northern feature gives the EM, for any assumed temperature, of the ionized gas along the line-of-sight \\citep{W00}.  The EM is a function of the electron density and the absorbing region thickness so its value constrains those parameters. The upper panel of Figure~\\ref{recomb10000} shows the constraints for an assumed temperature of $10^4$ K.  Constraints for a range of values of the EM are shown representing the range observed over different positions along the northern feature. Two other limits are also shown. The horizontal line is at 2.5 pc, the projected core distance of the centroid of the counterjet.  This is a rough upper limit for the thickness of the absorbing region if it is associated with the accretion disk.  For the observed emission measures, this thickness limit constrains the density to values above about $10^4$ cm$^{-3}$. A second limit, given by the lack of significant Thomson optical depth, is also shown.  Given the observed free-free absorption, the Thompson optical depth would only be important at very large sizes and low densities.  This could become an issue if the absorbing region is actually far enough from the AGN core to avoid severe line broadening by radiation damping.  If saturation and radiation damping were not a problem, the non-detection of H65$\\alpha$ could be used, along with the measured EM, to place some constraints on the physical conditions of the free-free absorbing gas.  For completeness, we will show those constraints.  The constraints would be of interest if the absorbing region is actually farther from the continuum radiation source than we believe or if for some reason the radiation damping is not as severe as we believe.  Using the measured EM and the assumed temperature, the optical depth of maser-like recombination line emission that would be expected from the absorbing gas can be calculated for a range of values of density and line width.  The lower panel of Figure~\\ref{recomb10000} shows the results of such calculations for an emission measure that is midway between the values obtained from the continuum absorption.  The calculations are based on departure coefficients very similar to those of \\citet{SB79} and are for the case when external radiation does not affect the level populations or line width.  Optical depths are plotted for several different line widths.  The figure also shows the upper limit to the line-to-continuum ratio of 0.15 that was derived above from the data.  In the absence of saturation or radiation damping, Figure~\\ref{recomb10000} shows that the recombination line non-detection would imply that the absorbing gas must have an appropriate combination of high density or high line width, presumably due to turbulence or other dynamic effects.  Or it could be at a velocity outside the observed window.  Note that the thermal line width for the assumed temperature is lower than the smallest velocity width shown.  To put the velocity range and line widths observed into some perspective, consider that the orbital velocity at about 3.3 pc from a $4\\times10^8 M_\\sun$ black hole is about 700 \\kms, or about half of the total observed velocity range.  That mass is clearly uncertain but is what was estimated by \\citet{WH01} using the correlation between black hole mass and bulge velocity dispersion.  The 3.3 pc distance is just the observed offset of 2.5 pc deprojected along the disk for the geometry described earlier.  That orbital velocity shows that it is possible for gravitationally bound material in the region of interest to be moving fast enough to produce line widths high enough to explain our non-detection.  But it also shows that, if the absorbing material is moving with the disk, it could be observable.  This would depend on the reasonable assumptions that the disk is moving transversely where it crosses the sight lines to the far-side jet, and that it has an internal velocity spread that is much smaller than the orbital velocity.  The Broad Line Region (BLR) line widths in NGC~1275 are of the order of our total frequency span \\citep[see, for example,][]{N99}.  A recombination line that wide would both not be sufficiently intense to observe, according to Figure~\\ref{recomb10000}, and would be washed out by our primary data reduction technique.  But typical BLR sizes, determined from reverberation mapping, are much smaller than the region probed by the sight lines through the free-free absorbing material.  Therefore the BLR line widths would only be a problem if the high velocities are maintained well beyond the region that emits the optical lines.   The most rudimentary models of accretion disks predict temperatures at distances such as those of interest here (few parsecs) that are too low to imply the presence of ionized gas.  But the free-free absorption shows that such gas is present.  Other lines of evidence, based on attempts to understand optical lines, also suggest that ionized gas is present \\citep[see][and references therin]{C99}.  It is possible that such gas is not in the plane of the disk, but rather in an atmosphere or wind above the disk that is ionized by energetic photons from the central regions of the system.  If this is the case, the constraints shown in Figure~\\ref{recomb10000} apply to the ionized atmosphere or wind, not the neutral core of the disk.  If a wind is involved \\citep[see, for example][]{K94}, it is likely that the the velocity would be outside the observed window, providing a natural explanation for the lack of an observed recombination line even if it weren't for the effects of the external radiation."
    },
    "0212/astro-ph0212520_arXiv.txt": {
        "abstract": "We review selected measurements of the galaxy luminosity function including the global field, the local group, the local sphere, nearby clusters (Virgo, Coma and Fornax) and clusters in general. We conclude that the overall cluster luminosity function is consistent with the global luminosity function over the magnitude range in common ($-22 \\leq M_B -5\\mbox{log}h_{0.68} \\leq -17 $). We find that only in the core regions of clusters ($r \\leq 300$ kpc) does the overall form of the luminosity function show significant variation. However when the luminosity function is subdivided by spectral type some further variations are seen. We argue that these results imply: substantial late infall, efficient star-formation suppression, and the confinement of mass-changing evolutionary processes to the core regions only. ",
        "introduction": "The Schechter luminosity function \\cite{Schechter1976} has been the standard expression for representing the space density of galaxies over the past 25 years.  It has strengths and weaknesses. The strengths are: its connection to the fundamental theory of the growth of initial mass perturbations \\cite{PS74}, its overall simplicity (with three free parameters: the characteristic luminosity $L^*$; the normalisation $\\phi^*$; and the faint-end parameter $\\alpha$), and its simple analytical connection to the luminosity-density (and ultimately via a mass-to-light ratio to the galaxian matter-densities).  Its weaknesses are: the correlation of the three parameters, the critical dependence of all three parameters upon the ``turn-over'' region, and its inability to reflect deviations of the space-density from a simple power-law at faint luminosities.  Over the past two decades numerous attempts have been made to constrain the three defining Schechter function parameters in both field and cluster environments producing a range of conclusions from a ubiquitous luminosity function to strong environmental dependencies. Here we showcase selected recent results from a variety of methodologies, to address this specific question as to whether the galaxy luminosity function is ubiquitous or not.  \\begin{figure} \\includegraphics[height=7.5cm,width=12.0cm]{spd_fig1.ps} \\caption{The recently published blue-band Schechter functions transformed to the standard Johnson B system. Surveys shown are: APM (Loveday et al 1992); ESP (Zucca et al 1997); SSRS2 (Marzke et al 1998); NOG (Marinoni et al 1999); SDSS1 (Blanton et al 2001); CS (Brown et al 2001); 2dFGRS (Norberg et al 2002); SDSS2 (Blanton et al 2003). } \\end{figure} ",
        "conclusions": "  \\noindent (1) Cluster infall is an ongoing process with a substantial fraction of the cluster population infalling in recent times. This explains the universality of the LF from the field to rich clusters and constrains any epoch of major merging to have occurred prior to the epoch of cluster formation.  \\noindent (2) Star-formation is effectively and efficiently inhibited in the infalling population however the halo merger rate is low. This explains the variation with spectral type but not in the overall LF. Hence while galaxies may have shifted their spectral classes their broad-band output remains mostly unchanged.  \\noindent (3) Dramatic evolutionary processes (merger, harassment etc) resulting in the construction of the cD/D galaxies and destruction of the dwarf population is confined to the core region. This explains the dips and upturns seen in Coma, the 2dFGRS core sample and numerous deep background subtracted core studies.  ~  \\noindent This is a data rich time for this field with the development of wide field mosaics, the Advanced Camera for Surveys and the upcoming large-format infrared detectors. We look forward to the power of these technologies being brought to bear on this unfolding story.   \\begin{table}[h] \\caption{A summary of Schechter LF estimates.} \\begin{tabular}{ccll} \\hline Survey     & Limit   & $M_{B}^{*}-5\\mbox{log}h_{0.68}$ & $\\alpha$ \\\\ \\hline {\\bf GLOBAL}         & $M_B \\leq -17$ & $-20.47 \\pm 0.2$ & $-1.17 \\pm 0.02$ \\\\ {\\bf LOCAL GROUP}    & $M_B \\leq  -7$ & $-21^{+4}_{-2}$        & $-1.2^{+0.24}_{-0.14}$ \\\\ {\\bf LOCAL SPHERE}   & $M_B \\leq -11$ & $-20.1^{+1.2}_{-0.6}$  & $-1.19^{+0.05}_{-0.06}$ \\\\ {\\bf VIRGO}          & $M_B \\leq -11$ & $-20.2^{+0.7}_{-0.5}$  & $-1.32^{+0.03}_{-0.03}$ \\\\ {\\bf FORNAX}         & $M_B \\leq -12$ & $-20.2^{+0.5}_{-0.4}$  & $-1.23^{+0.04}_{-0.03}$ \\\\ {\\bf COMA}           & $M_B \\leq -14$ & $-20.1^{+0.25}_{-0.20}$  & $-1.01^{+0.04}_{-0.05}$ \\\\ {\\bf Garilli et al}$^1$  & $M_{B} \\leq -17$ & $-20.6 \\pm 0.1$  & $-0.97^{+0.04}_{-0.02}$ \\\\ {\\bf Paolillo et al}$^1$ & $M_{B} \\leq -16$ & $-20.5^{+0.13}_{-0.17}$ & $-1.07^{+0.09}_{-0.07}$ \\\\ {\\bf Goto et al}$^1$     & $M_{B} \\leq -16.5$ & $-21.47 \\pm 0.11$ & $-1.00 \\pm 0.06$ \\\\ {\\bf 2dFGRS FIELD}   & $M_B \\leq -17$ & $-20.35 \\pm 0.04$  & $-1.19 \\pm 0.01$ \\\\ {\\bf 2dFGRS CLUSTER} & $M_B \\leq -17$ & $-20.62 \\pm 0.07$  & $-1.28 \\pm 0.03$ \\\\ \\hline \\end{tabular}  \\noindent $^1$ $g$-band data converted to $B$ assuming $B=g+0.54$  \\vspace{-1.0cm}  \\end{table}  \\vspace{-0.3cm}"
    },
    "0212/astro-ph0212185_arXiv.txt": {
        "abstract": "We argue that tracing star formation histories with GAIA using main sequence turn-off (MSTO) point dating will mainly be effective in cases of very mild interstellar extinction ($E_{B-V}<0.5$). For higher reddenings the MSTO approach will be severely limited both in terms of traceable ages ($t<0.5$ Gyr at 8 kpc; $E_{B-V}=1.0$) and/or distances ($d=2$ kpc if $t\\leq15$ Gyr; $E_{B-V}=1.0$), since the MSTO will be located at magnitudes too faint for GAIA. AGB stars may alternatively provide precise population ages with GAIA for a wide range of ages and metallicities, with traceable distances of up to 250 kpc at $E_{B-V}=0$ (15 kpc if $E_{B-V}=2.0$). It is essential however that effective temperatures, metallicities, and reddenings of individual stars are derived with the precision of $\\sigma(\\log T_{\\rm eff})\\sim0.01$, $\\sigma([M/H])\\sim0.2$, and $\\sigma(E_{B-V})\\sim0.03$, to obtain $\\sigma(\\log t)\\sim0.15$. This task is quite challenging for GAIA photometry and spectroscopy, though preliminary tests show that comparable precisions may be achieved with GAIA medium band photometry. ",
        "introduction": "\\begin{figure}[t] \\plotone{kucinskas_fig1.eps} \\caption{Left: maximum age for which MSTO point can be dated at 8 kpc with GAIA, for different [Fe/H] and $E_{B-V}$. Right: maximum distance at which stellar population up to 15 Gyr old can be dated using MSTO with GAIA, for different [Fe/H] and $E_{B-V}$.} \\end{figure}   Out of the numerous methods available for assessing ages of stellar populations only a few will be practically applicable for use with GAIA, because of obvious limitations related to the detection limits of GAIA etc.\\ (e.g., Cacciari 2002). Though the main sequence turn-off (MSTO) point approach is perhaps the most reliable and accurate of all dating methods available today, it will be severely distance-limited when used with GAIA.  The extent of these limitations is illustrated in Fig.~1. The left panel shows the maximum age obtainable with GAIA employing MSTO fitting for stellar population placed at a distance of 8 kpc (assuming that the MSTO point is located at $V\\sim18.5$, while the detection limit of GAIA is $V=20$). The right panel shows the maximum distance at which ages up to 15 Gyr may be quantified (in both cases, the MSTO point on the isochrones of Girardi et al.\\ (2000) was dated). It is obvious that in cases of negligible interstellar extinction the MSTO approach may work very efficiently with GAIA, providing reliable ages of stellar populations up to $\\sim8$ kpc. However, even mild interstellar extinction changes the situation dramatically, shifting the MSTO point below the detection limits of GAIA for many stellar populations within the Galaxy. ",
        "conclusions": ""
    },
    "0212/astro-ph0212466_arXiv.txt": {
        "abstract": "The normal parameters are a non--linear transformation of the cosmological parameters whose likelihood function is very well--approximated by a normal distribution. This transformation serves as an extreme form of data compression allowing for practically instantaneous calculation of the likelihood of any given model, as long as the model is in the parameter space originally considered.  The compression makes all the information about cosmological parameter constraints from a given set of experiments available in a useable manner.  Here we explicitly define the normal parameters that work for the current CMB data, and give their mean and covariance matrix which best fit the likelihood function calculated by the Monte Carlo Markov Chain method.  Along with standard parameter estimation results, we propose that future CMB parameter analyses define normal parameters and quote their mean and covariance matrix. ",
        "introduction": "The challenge of turning a CMB dataset into cosmological parameter constraints is one that has been solved by a series of data compressions.  The information in the time--ordered data is compressed into a map \\citep{wright96,tegmark97a}.  The information in the map is then compressed into a power spectrum \\citep{tegmark97b,bond98,bond00,wandelt01,bartlett00,hivon02}. Finally, the information in the power spectrum is then compressed into cosmological parameters \\citep{lineweaver97,benoit02}.  This last step, however, is more of a data ``explosion'' than a data compression. Although the number of parameters is indeed small ($\\sim 10$) the non--Gaussian distribution of their errors may be characterized by $10^{10}$ numbers for grid--based likelihood evaluation \\citep{tegmark00} or perhaps as little as $30,000$ for the Monte-Carlo Markov Chain method \\citep{christensen01}.  This explosion has undesirable consequences.  The full constraining power of the data is not used by the cosmological community.  Typically authors report not the full (cumbersome) likelihood, but projections or marginalizations of it down to one or at most two--dimensional spaces. Given the near--degeneracies that exist in the probability distribution, and its non-Gaussian nature, these final steps lose significant information.  Our work here provides the final step of compression in the data analysis pipeline.  We compress the probability distribution of the cosmological parameters down to the $\\sim 50$ parameters of an analytic form for the likelihood (a mean and covariance matrix). This step makes the full information in the parameter likelihood function easily useable.  Our procedure is analogous to the `radical compression'\\citep{bond00} used for compressing the uncertainty in the power spectrum estimates, whose distribtuion is also non--Gaussian.  In both cases a non--linear variable transformation is used, with the property that the transformed variables are well--approximated by a normal distribution.  We were inspired to search for a set of normal parameters by a proposal that the $C_l$ can have a linear dependence on a set of well-chosen parameters that span the space of possible models \\citep{kosowsky02}.  Clearly such a parameter set would be tremendously valuable for parameter estimation since the $C_l$ for any given model could be calculated practically instantaneously simply by summing the terms in the first order Taylor expansion. Such a set would also, given sufficiently Gaussian errors on $C_l$,  have errors with a normal distribution.  Although inspired by \\citet{kosowsky02}, our goal here is entirely different.  It is to find a set of parameters (the normal parameters) for which the likelihood is Gaussian.  Although the approximate linearity demonstrated by \\citet{kosowsky02} gave us confidence to try to find a set of normal parameters, it by no means guaranteed success. The linear-response approximation has not yet been demonstrated to be sufficiently accurate for purposes of parameter estimation given MAP data (and certainly not given current data) and the errors on $C_l$ are not actually Gaussian.  While \\citet{kosowsky02} are developing a tool for parameter estimation from CMB data, our work is aimed at being able to easily take full advantage of such parameter estimation.  The most straightforward application of CMB data compressed in this form is for simultaneous analysis of CMB constraints with those from other cosmological probes.  In the past to do this analysis one would have to entirely reproduce the reduction from bandpowers (plus their window functions, offset log--normal parameters, Fisher matrices, and calibration uncertainties) to cosmological parameter likelihoods. This procedure requires the organization of a lot of data as well as hundreds of thousands of angular power spectrum calculations.  In contrast, with the compression to normal parameters in hand one need only perform a variable transformation and evaluate a $~10$--dimensional Gaussian.  We use a similar set of parameters as KMJ, and demonstrate explicitly the validity of the normal approximation by numerical computation of the likelihood given current CMB data.  One can also simplify a probability distribution by choosing linear combinations of the parameters that diagonalize the covariance matrix, the so--called parameter eigenmodes \\citep{efstathiou99}.  However, the resulting parameters are still dependent and non--Gaussion if the distribution of the original parameters is non-Gaussian. We do find diagonalization of the covariance matrix {\\em of the normal parameters} to be useful, both for our fitting procedure and for gaining an understanding of the nature of the constraints the data places on the parameter space.  We test our procedure on current data and provide the mean and covariance matrix to the normal parameters that best fit the likelihood given current data.  As we go to press, this fit is now out of date due to the data from the Wilkinson Microwave Anisotropy Probe ({\\it WMAP}) \\footnote{http://map.gsfc.nasa.gov}. However, we expect these data to further improve the validity of the normal approximation to the likelihood and thus our method will be of lasting value.  In section II we describe the normal parameters. In section III we describe our calculation of the exact likelihood via the Monte Carlo Markov Chain method \\citep{christensen01}.  In section IV we describe our procedure for finding the mean and covariance matrix that provide the best fit to our MCMC--calculated likelihood. In section V we present our results. In section VI we discuss applications and conclude. ",
        "conclusions": "We have defined a set of parameters that, given current data, have a likelihood that is normally distributed.  Of course, our simple fit to the likelihood surface will very soon be out-of-date due to the expected MAP data.  Still, we expect the technique to continue to be useful.  Generally, as data become more constraining, the Gaussian approximation becomes better.  The challenge in the future will be to provide good parameterizations for those parameters that remain poorly determined. Some of these only have their effects at low $\\ell$, such as the parameters governing reionization.  For these it may be useful to define parameters for which $\\ln{C_l}$ is linear, since the errors in this quantity are very nearly normally distributed at low $\\ell$ \\cite{bond00}.  Other physical effects, such as gravitational lensing, are only important at high $\\ell$ where, if instrument noise is the dominant contribution to error in $C_l$, then $C_l$ will be nearly normally distributed.  For these one would want to use parameterizations for which $C_l$ is linear.  An analytic form for the likelihood makes the actual application of CMB parameter constraints to cosmological questions much easier. One application is to quickly calculate constraints on various combinations of the cosmological parameters such as $\\sigma_8$ \\citep{holder02} or the age, $t_0$ \\citep{knox01}.  Another is to combine the CMB results with those from other probes to improve constraints or search for inconsistencies \\citep{wang01}. Or one can forecast expected constraints from a combination of the CMB and other probes such as supernovae \\citep{frieman02} or galaxy cluster number counts \\citep{levine02}.  Although these applications are possible without an analytic form for the CMB likelihood, such a form greatly simplifies the analysis.  Use of normal parameters will also improve the efficiency of Monte Carlo evaluations of the likelihood, by providing an analytic generating function that is well--matched to the actual likelihood. Although the statistical properties of the post--convergence chain do not depend on the generating function used, the time required for convergence depends critically on the generating function. In particular the chain rapidly converges if the generating function approximates the likelihood well. Clearly, this issue becomes more important for larger parameter spaces. A practical way to implement the use of normal parameters in generating Monte Carlo Markov chains would be to create a small chain, derive the normal parameters and the gaussian approximation to the likelihood, and then use this gaussian likelihood as the generating function.  In this paper we have explicitly given an analytic form for the likelihood (given a certain data set) of cosmological parameters and shown it to be an excellent approximation to the MCMC--calculated likelihood. We view this as the final step in a long data-reduction chain that begins with the time--ordered data."
    },
    "0212/hep-ph0212365_arXiv.txt": {
        "abstract": " ",
        "introduction": "Most aspects of high energy physics beyond the standard model can only be tested by going to very high energies, which are by far greater than those accessible by present, or even future, terrestrial accelerators.  Cosmology has offered a way to {\\sl experimentally} test new theories of fundamental forces. In these lectures, I will shortly present the fascinating interplay between particle physics and cosmology, as provided by topological defects.  \\par Many particle physics models of matter admit solutions which correspond to a class of topological defects, that are either stable or long-lived.  Provided our understanding about unification of forces and the big bang cosmology are correct, it is natural to expect that such topological defects could have formed naturally during phase transitions followed by spontaneously broken symmetries, in the early stages of the evolution of the universe.  Certain types of defects lead to disastrous consequences for cosmology and thus, they are undesired, while others may play a useful r\\^ole.  \\par Spontaneous symmetry breaking is an old idea, having its origin in the field of condensed matter physics, where it is described in terms of the order parameter, while in particle physics, symmetry breaking is described in terms of a scalar field, the Higgs field. The symmetry is called spontaneously broken (SSB) if the ground state is not invariant under the full symmetry of the Lagrangian density. Thus, the vacuum expectation value of the Higgs field is non-zero. In quantum field theories, broken symmetries are restored at high enough temperatures, as one can easily see from the finite temperature effective potential, which can be calculated from the zero temperature classical potential in loop expansion.  \\par In three spatial dimensions, four different kinds of topological defects can arise. Whether or not topological defects will develop during a symmetry breaking phase transition as well as the determination of their type both depend on the topology of the vacuum manifold ${\\cal M}$. The properties of ${\\cal M}$ are usually described by the $n^{\\rm th}$ homotopy group $\\pi_n({\\cal M})$ which classifies dinstinct mappings from the $n$-dimensional sphere $S^n$ into the manifold ${\\cal M}$.  If ${\\cal M}$ has disconnected components, or equivalently if the order $n$ of the non-trivial homotopy group is $n=0$, then two-dimensional defects, called {\\sl domain walls}, form.  The spacetime dimension $d$ of the defects is given in terms of the order of the non-trivial homotopy group by $d=4-1-n$. If ${\\cal M}$ is not simply connected, in other words if ${\\cal M}$ contains loops which cannot be continuously shrunk into a point, then {\\sl cosmic strings} form. A necessary, but not sufficient, condition for the existence of stable strings is that the first homotopy group (the fundamental group) $\\pi_1({\\cal M})$ of ${\\cal M}$, must be non-trivial, or multiply connected. Cosmic strings are line-like defects, $d=2$. If ${\\cal M}$ contains unshrinkable surfaces, then {\\sl monopoles} form, for which $n=2, ~d=1$. Finally, if ${\\cal M}$ contains non-contractible three-spheres, then {\\sl textures} form. Textures are event-like defects with $n=3, ~d=0$.  \\par Topological defects are called local or global. The energy of local defects is strongly confined, while the gradient energy of global defects is spread out over the causal horizon at defect formation.  \\par In these lectures I will basically concentrate on cosmic strings, which have very intriguing properties and interesting cosmological implications~\\cite{review-td,hk}. I will discuss the r\\^ole which they could play in the large-scale structure and the anisotropies of the cosmic microwave background, as well as the r\\^ ole of cosmic strings in explaining various high energy phenomena.  I will then summarize the gravitational effects of cosmic strings networks.  Finally, I will discuss topological defects within brane world cosmology.  \\par This review is organized as follows: In Section 2, I describe cosmic strings in the context of field theory. I first discuss two simple models, the abelian-Higgs model and the Goldstone model, which lead to local and global strings respectively. I then summarize string formation and string dynamics.  After this short introduction to the physics of cosmic strings, I proceed with their physical/astrophysical consequences.  In Section 3, I analyze the mechanism of cosmic structure formation with topological defects. I first describe the basic observable of the cosmic microwave background, and the physical mechanisms which are responsible for the microwave background anisotropies. I compare topological defects models versus inflationary ones, as the two alternative mechanisms for structure formation within gravitational instability theory. Concluding that topological defects are excluded as the main source of the large-scale structure, I address the question of finding the more {\\sl natural} class of inflationary models, meaning inflationary scenarios motivated by high energy physics. I then discuss the implications of these findings for particle physics models. In Section 4, I describe how cosmic strings could explain various high energy phenomena (baryon asymmetry, ultra-high energy cosmic rays, and gamma ray bursts). In Section 5, I discuss the gravitational effects of cosmic strings networks.  In Section 6, I briefly comment on the r\\^ole of topological defects in brane world cosmology. I close this review with the conclusions in Section 7. ",
        "conclusions": "Topological defects in general, and cosmic strings in particular, present a very active subject of modern cosmology. These objects arise in a wide class of elementary paricle physics models and they are inevitably formed during phase transitions in the early universe. They  present a fruitful interplay between cosmology and high energy physics. Lately defects gained even more interest since one can {\\sl study} various defects aspects in condensed matter experiments.  \\par In these lectures I briefly presented topological defects from the point of view of field theories and I then discussed some of their cosmological/astrophysical consequences.  \\par Topological defects are ruled out as the main mechanism for the observed large-scale structure and for the anisotropies of the cosmic microwave backround radiation. Nevertheless, cosmic strings can still play a non-negligible r\\^ ole in this aspect.  The task we are facing now is to use this knowledge (combining the predictions of the models with defects together with the observational data) to choose an acceptable inflationary scenario, and also to constrain particle physics models. In addition, topological defects may have played a r\\^ ole in some other physical/astrophysical questions which I have briefly discussed here."
    },
    "0212/hep-ph0212223_arXiv.txt": {
        "abstract": " ",
        "introduction": "The fact that the observed arrival directions of UHECR [1,7,8] show significant clustering [7,12,9] suggests that their sources may be identified using already existing data.  Poor (in astronomical standards) angular resolution of the UHECR experiments and bending of particle trajectories by Galactic and extragalactic magnetic fields does not allow an event-by-event identification. One therefore has to relay on statistical methods, i.g. on calculation of the angular correlation function [9,10]. In this approach, the significance of correlations at a given angular scale $\\delta$ may be characterized by the probability $p(\\delta)$ that the observed excess of events around sources at a scale $\\delta$ may have occurred by chance. When $p(\\delta)$ is small, the correlation is physical as it cannot be explained by a statistical fluctuation.  Apart from the angle, $p(\\delta)$ depends on the set of potential sources and other parameters like assumed charges of particles and parameters of the Galactic magnetic field (GMF). Significance of correlations may be used to determine the physical values of these parameters and the subset of actual sources by looking at the minima of $p(\\delta)$.  The choice of candidate sources  depends on the assumptions about the nature of UHECR. In the most conservative approach sources are associated with astrophysical accelerators, among which active galactic nuclei (AGN) may be singled out for their total energetics and good conditions for particle acceleration [2,3]. Even in such powerful accelerators, the acceleration of protons to energies of order $10^{20}$~eV requires exceptional conditions. It is therefore quite likely that only a small fraction of AGNs can produce UHECR.  Comparing the number of isolated cosmic ray events to the number of close pairs and higher multiplicity clusters, one may estimate the number of point sources to be of order several hundred [4]. This agrees with the expectation that only a special class of AGNs can be UHECR sources. A subclass of ANGs, blazars, are thought to have relativistic jets directed along the line of sight. Therefore, if UHECR do correlate with AGNs, they should correlate with blazars. A subset of blazars called BL Lacertae objects (BL Lacs) is characterized, in addition, by narrow (or absent) emission lines. This latter property suggests low ambient matter and radiation densities near the central engine and may provide favorable conditions for particle acceleration to highest energies. In addition, there are indications of negative evolution of BL Lacs (there were less of them in the past), which makes the GZK cut-off less sharp. We conclude that, on theoretical grounds, BL Lacs are potentially promising sites of UHECR acceleration and should be selected as a set of candidate sources for correlation analysis.  The most recent catalog of AGN contains 350 confirmed BL Lacs [13], the closest being at $z\\simeq 0.03$.  Our analysis shows significant correlations between UHECR and most powerful BL Lacs [10]. The correlations occur at angles comparable to the experimental angular resolution. It implies that BL Lacs are sources of UHECR.  Can one find, in the existing data, independent signatures supporting this result? This is the main question which we address in the present talk. We argue that these signatures do exist. Moreover, they are compatible with the assumption that UHECR of moderate energies are protons.  The correlation between UHECR and BL Lacs implies that the origin of (at least a part of) UHECR is acceleration. This has a number of consequences. First, one should expect that protons constitute at least part of the UHECR flux. Since protons are charged, they are deflected in the Galactic and extragalactic magnetic fields. The correction for these deflections should improve correlations with sources. This effect is indeed present in the data [11]. Second, any acceleration process is accompanied by energy losses. A substantial part of this energy goes into electromagnetic channel and ends up in the EGRET energy range [2]. Thus, one may expect the connection between sources of UHECR and gamma radiation. There are strong arguments in favor of this connection [5].  Finally, any {\\it chance coincidences} between cosmic rays and BL Lacs should be distributed over the sky randomly (i.e. reflecting only the local density of BL Lacs and exposure of cosmic ray experiments).  Thus, any significant deviation in the distribution of {\\it correlating} rays over the sky from this expectation speaks in favor of real physical connection between cosmic rays and BL Lacs. We show that correlating rays {\\it are not} distributed randomly.  In three subsequent sections we discuss these signatures in more detail. The last section contains discussion and conclusions. ",
        "conclusions": ""
    },
    "0212/astro-ph0212136_arXiv.txt": {
        "abstract": "GAIA data will create a precise 3-dimensional map of  the Galaxy, providing positional information, radial velocities, luminosity, temperature and chemical composition of a representative sample of stars. Here we present a new implementation of the Padova Galaxy Model, where kinematics simulations are included. A few examples of application to GAIA science are discussed, in particular concerning radial velocities determinations of the thin and thick disk. ",
        "introduction": "GAIA  will provide detailed phase space coordinates for about 1 billion stars within a sphere of 20 Kpc. In addition to this  on board multicolour photometry and spectroscopy will give complete chemical measurements (including $\\alpha$, $s$, $r$ process elements) down to V=17-19 mag. Owing to the precision of the data, the problem of the formation of the Milky Way will be addressed. In the standard cold dark matter scenario the Galaxy should have formed by  merges of small substructures. However this scenario is at odds with several observational constraints: i.e. a  small scale length of the disk of the order of 300 pc is expected instead of 2000-3000 pc  (Steinmetz \\& Navarro 1999); the formation process predicts that the Milky Way should have a thousand of satellites which are missing in the Local Group. A substantial revision of the model is necessary, by means of a detailed comparison with high quality data. GAIA will bring into evidence fossil remnants of the galaxy formation tracing back the star formation history and revealing gradients in the star formation intensity, chemical abundances, and kinematics across the disk. Chemical and age gradients are not expected inside the thin disk. N-body simulations show that radial mixing can wash out gradients close to the Galactic plane after a few Gyr. Concerning the thick disk, if it was formed by an heating event (merge with a satellite) traces of it should still be detectable (Freeman \\& Hawthorn 2002). In fact vertical gradients in  metallicity and velocity dispersions which are still uncertain from the present data  are possibly not washed out by the radial mixing far from the Galactic plane and can be used to trace back the formation history. In order to simulate Galactic data we update the Padova Galaxy Model including velocity  space and we apply it to the GAIA science.  \\section {Padova Galaxy Model}  The Galaxy is modeled with the code already described  by Bertelli et al. (1995) and revised as  in Vallenari et al (2000) where more detail can be found. The Padova model  has been newly updated including:  1) the usage of new stellar tracks from Z=0.0001 till Z=0.03 with low mass stars down to 0.15 (Girardi et al 2000). 0.1 M$_\\odot$ track is taken from  Baraffe et al (1998);  2) the use  of the carbon star models taking into account the effect of the variation of the stellar molecular  opacities during the evolution as a result of the dredge-up (Marigo 2002);  3) the extinction along the line of sight derived following   Drimmel \\& Spergel (2001) model obtained from COBE-DIRBE infrared data.  4) the velocity space  simulated in a consistent way as described in the following Section.  \\section {The kinematic model}  The velocity distribution for the whole disk has been computed using the velocity ellipsoids formulation by Schwarzschild. Concerning the thin disk, the assumed local values of velocity dispersions  are taken as in Mendez et al (2000). The V$_{lag}$ is derived following Binney \\& Tremaine (1994). The diagonal terms of the dispersion velocity tensor are from Lewis \\& Freeman (1989):  \\begin{eqnarray} \\sigma^2_{RR} &=& \\sigma^2_{RR,0}exp(-(R-R_0)/H_R)\\\\ \\sigma^2_{\\Phi \\Phi} &=& 1/2 (1+(dln V_{LSR} (R))/(dlnR)) \\times \\sigma^2_{RR})\\\\ \\sigma^2_{ZZ}&=&\\sigma^2_{ZZ,0}exp(-(R-R_0)/H_{R}) \\end{eqnarray}  The vertical gradient in the velocity dispersions is taken from Fuchs \\& Wielen (1987) for $\\sigma^2_{RR}$ and $\\sigma^2_{\\Phi \\Phi}$ and the non vertical isothermality of the thin disk from Amendt \\& Cutterford (1991) which is in good agreement with the observations till 1 Kpc for $\\sigma^2_{ZZ}$. The off diagonal term  $\\sigma_{RZ}$ is derived from Amend \\& Cutterford (1991):  \\begin{equation} \\sigma_{RZ}^2(R,Z)=\\sigma^2_{RZ}(R,0)+z\\partial/{\\partial z} \\sigma^2_{R,Z}(R,0) \\end{equation}  where  \\begin{equation} \\partial/{\\partial z} \\sigma^2_{RZ}(R,0)= \\lambda(R)({\\sigma^2_{RR}-\\sigma^2_{ZZ}}/R)(R,0) \\end{equation}  and  \\begin{equation} \\lambda(R)=(R^2 \\Phi_{Rzz}/{3\\Phi_{R}+R \\Phi_{RR}-4R\\Phi_{ZZ}})(R,0) \\end{equation}  where $\\Phi_{RZZ}, \\Phi_{RR},\\Phi_{ZZ}$ are the derivatives of the Galactic potential $\\Phi$  obtained from the density model by Dehnen \\& Binney (1998) by inverting the Poisson equation with the Bessel integrals (Quinn \\& Goodman 1986). $\\lambda$ is an approximate expression  of  the vertical tilt of the velocity ellipsoid close to the Galactic  plane (at z=0).  $\\lambda$ is 1  in the case of a spherical potential, when the ellipsoid is pointing towards the Galactic center , while  $\\lambda$ is 0 for a cylindrical potential when the ellipsoid is always parallel to the Galactic  plane. Amendt \\& Cutterford (1991) show that $\\lambda$ is related to the mass gradient in  the Galactic plane. No vertex deviation is included in the model although a deviation decreasing from 25$^0$ for young stars to near zero for an old population has been found by various authors (Dehnen \\& Binney 1998, Bienayme 1999, Soubiran et al 2002). The thick disk is  isothermal with $(\\sigma_{RR}, \\sigma_{\\Phi \\Phi}, \\sigma_{zz})=(70,50,45)$ Km s$^{-1}$  as starting values. The $V_{lag}$ is assumed to have a canonical value of 35  Km s$^{-1}$. The possibility of simulating a vertical gradient in the rotational velocity of the thick disk is included as suggested by Chiba \\& Beers (2000). The halo  has $(\\sigma_{RR}, \\sigma_{\\Phi \\Phi}, \\sigma_{zz})=(130,95,95)$ Km s$^{-1}$. The projection matrixes of the space velocities are derived from Mendez et al (2000). In the following Sections we show a few examples of applications to the GAIA science. In particular the thin and thick disk are discussed. ",
        "conclusions": ""
    },
    "0212/astro-ph0212070_arXiv.txt": {
        "abstract": "We report a $\\simeq 3 \\sigma$ detection of a 21.82 \\AA \\/ absorption feature in the direction of the BL Lac MRK~421 which is interpreted as O~{\\sc vii} K$\\alpha$ at $z\\simeq 0.01$. This corresponds to the redshift of a H absorber in a cosmic void detected by HST along the line of sight of the same object. The 21.82 \\AA  \\/ line proves the existence of warm/hot intergalactic medium (WHIM) outside the Local Group, in agreement with current models of structure formation in the Universe. The WHIM at $z\\sim 0.01$ has temperature of $\\lsimeq (2.5-4) \\times 10^6$~K for gas densities in the range $10^{-6}-1$ atom cm$^{-3}$. We also detect  O~{\\sc vii},  O~{\\sc viii} and  Ne~{\\sc ix} absorption lines with zero velocity, possibly connected with WHIM in the Local Group. ",
        "introduction": "In the local Universe $\\sim 50 -60$\\% of the baryons visible at $z>2$ and predicted by the Standard Big-Bang nucleosynthesis are undetected \\citep[e.g.][]{fuk98}. Numerical simulations predict that such baryons, in the form of Ly$\\alpha$ clouds (i.e. gas at $T <10^4$ K) and visible in the UV spectra of bright AGNs, have been shock-heated during the formation of the structures in the Universe to temperatures of $\\sim 10^5-10^7$~K (i.e. the warm hot intergalactic medium, WHIM). The most efficient way to detect such gas is through resonant absorption lines from highly ionized metals (e.g. O{\\sc vi},O{\\sc vii},O{\\sc viii},Ne{\\sc ix}) in the far UV and soft X-ray spectra of background sources \\citep[e.g.][]{hel98,fan00}. Current far UV observations proved the existence of the low temperature tail ($T \\sim (1-5) \\times 10^5$ K) of the WHIM through the detection of O{\\sc vi} up to $z\\sim 0.2$ \\citep[e.g.][]{sem00, tri01}. However the bulk of the WHIM baryons \\citep[$\\sim 70$\\%, e.g.][]{fuk98} are expected to lie at higher temperatures ($T\\sim 10^6- 10^7$~K) and their features should be detectable in the soft X-ray regime. {\\em Chandra} and {\\em XMM-Newton} high resolution spectroscopy ($R \\sim 100-500$) brought to the detection of a ``zero redshift'' O{\\sc vii} absorption feature at 21.6 \\AA \\/ in the spectra of four bright AGNs  (PKS~2155-304, Nicastro et al. 2002; Fang et al. 2002; Cagnoni et al. 2003;  3C~273, Fang, Sembach \\& Canizares 2003; H1821$+$643, Mathur, Weimberg \\& Chen 2002 and MRK~421, Nicastro et al. 2001), interpreted as the signature of WHIM present within the Local Group. The only published evidence of WHIM outside the Local Group is at $z\\sim 0.05$ \\citep{fan02}, however this feature is not confirmed in other {\\em Chandra}  and {\\em XMM-Newton} observations \\citep[e.g.][]{nic02,cag03}. In this paper I present  strong evidence for an  O{\\sc vii} K$\\alpha$ absorption line at  $z\\sim 0.01$ in the spectrum of the BL Lac object MRK~421. \\\\ The paper is organized as follows: \\S~\\ref{sec_obs} reports the {\\em XMM-Newton} observations of MRK~421 and  describes the data reduction; \\S~\\ref{sec_whim} contains the spectral fits and the discussion on the observed WHIM absorption features. \\S~\\ref{sec_sum} is a conclusive section containing a summary of the results. ",
        "conclusions": "\\label{sec_sum} I report  strong evidence of WHIM outside our local group of galaxies through the detection of a 21.82 \\AA  \\/ O~{\\sc vii} K$\\alpha$ absorption line in the   {\\em XMM-Newton} RGS-1 spectrum of MRK~421. The line redshift is $\\sim 0.01$,  in agreement with that derived for the strong H Ly$\\alpha$ absorption line detected by HST and related to an absorber in a cosmic void. The three closest galaxies are at 2.14, 3.98 and 4.12 $h_{70}^{-1}$~Mpc \\citep{pen00}. This detection is a firm proof the existence of a WHIM filament at $z\\sim 0.01$ possibly connecting the nearby galaxies with other denser regions of the sky, as expected by the models of the cosmic structures formation and evolution. I derive an upper limit on the gas temperature of $\\sim 2.5-4 \\times 10^6$~K for a gas density of $10^{-6} - 1$ atom cm$^{-3}$. The only known similar evidences are  a  $4.5 \\sigma$ detection of O~{\\sc viii} at $z\\sim 0.05$ reported by \\citet{fan02} in the {\\em Chandra} spectrum of PKS~2155-304 and the $\\sim 2 \\sigma$ detections of two O~{\\sc vii} lines at $z\\sim 0.2$ in the {\\em Chandra} spectrum of H1821$+$643 \\citep{mat02}. However \\citet{fan02} feature is not confirmed in other {\\em Chandra}  and {\\em XMM-Newton} observations \\citep[e.g.][]{nic02,cag03}.\\\\ I also detect the zero redshift absorption features of  O~{\\sc vii}, O{\\sc vii} and Ne{\\sc ix} with positions and EWs consistent with previous detections in the direction of other bright AGNs \\citep[e.g.][]{nic02,cag03}.\\\\"
    },
    "0212/astro-ph0212252_arXiv.txt": {
        "abstract": "We present wide field CCD photometry of a galactic globular cluster M92 obtained in the V and I bands with the CFH12K mosaic CCD at the Canada-France-Hawaii Telescope. A well-defined color-magnitude diagram is derived down to 5 magnitudes fainter than the cluster main sequence turn-off. After removing the background contribution, we obtain luminosity and mass functions, surface density profiles, and the surface number density maps of the stars belonging to the cluster. The surface density profile of all stars shows that the cluster's halo extends at least out to $\\sim30'$ from the cluster center in agreement with previous study, but the profile of faint stars at the very outer region of the cluster shows a different gradient compared with that of bright stars. For a mass function of the form $\\Phi(M) \\propto M^{-(1+x)}$, we find that the inner region ($5' < r < 9'$) of the cluster has $x\\simeq1.2\\pm0.2$, whereas the outer region ($9' < r < 15'$) has $x\\simeq1.8\\pm0.3$, clearly indicating a mass segregation of the cluster. An estimate of the photometric mass of the cluster implies that the remnant populations (white dwarfs and neutron stars) contribute at least 25\\% of the total cluster mass. The surface density map of M92 shows some evidence that the tidal tail of M92 may be oriented perpendicular to the direction toward the Galactic center. ",
        "introduction": "Globular clusters are good laboratories for testing theories of stellar dynamics and galactic structure. All globular clusters ultimately decay, through a combination of internal (two body relaxation, and stellar evolution) and external (gravitational shocks due to the Galactic bulge and disk, and dynamical friction) processes (Aguilar, Hut, \\& Ostriker 1988). Lee \\& Goodman (1995) and Gnedin \\& Ostriker (1997) predicted that possibly as many as half of the present-day Galactic globular clusters would be disrupted in another Hubble time. Many previous studies have shown that tidal shocks play a very important role in cluster evolution by accelerating the destruction of the cluster (e.g. Ostriker, Spitzer, \\& Chevalier 1972; Gnedin \\& Ostriker 1997; Gnedin, Lee, \\& Ostriker 1999). Lee \\& Ostriker (1987) predicted that, for a given tidal field, the outer structure of a globular cluster will be more populated than that of a King model, since the density at the tidal radius would not be exactly zero. In addition, the cluster is expected to have tidal tails because the equipotential surface is not spherical. The presence of tidal tails from globular clusters has begun to be explored only recently because of difficulties in obtaining accurate wide field photometry. The development of computing power and wide field digital detectors now allows us to study these very interesting parts of globular clusters in detail. Indications of tidal tails around globular clusters were found in several previous studies of globular cluster radial density profiles. Grillmair et al. (1995) examined the outer structure of 12 Galactic globular clusters, and found that 10 of their sample clusters showed extra-tidal wings in their surface density profiles. Later, Lehmann \\& Scholz (1997) reported the discovery of tidal tails in 5 out of 7 globular clusters. Evidence of tidal extensions are even found in four globular clusters in M31 (Grillmair 1996; Holland, Fahlman, \\& Richer 1997). Leon, Meylan, \\& Combes (2000) investigated 20 Galactic globular clusters with Schmidt plates and films, and found that all of the clusters that do not suffer from strong observational biases, show tidal tails. On the theoretical side, Combes, Meylan, \\& Leon (1999) used N-body simulations to show that two giant tidal tails are expected to exist along the galactic orbit of a globular cluster. Yim \\& Lee (2002) also examined the general evolution of the clusters using N-body simulation and showed that the cluster density profiles appear to become somewhat shallower just outside the tidal boundary, and that the directions of the tidal tails depend on the location in the Galaxy as well as the cluster orbit.  Recently, Odenkirchen et al. (2001) reported the discovery of two well defined tidal tails emerging from the sparse, remote globular cluster Palomar 5. They used the CCD images from Sloan Digital Sky Survey (SDSS). Except for this study (and those of the clusters in M31), all the previous observational works on globular cluster tidal tails were based on data from photographic plates. However, such data generally are not deep enough to examine the low mass stars which dominate the outer parts of globular clusters, and the photographic photometry tends to show greater scatter than that based on CCD data.  In the specific case of M92, Testa et al. (2000) investigated the dynamical structure using plates from the Digitized Second Palomar Sky Survey (DPOSS). Although they found some evidence for an elongation in the extra-tidal extension from their surface density map, the significance of their result was low. In this paper, we report the results of a similar analysis using deep CCD images. We surveyed a $2^\\circ \\times 2^\\circ$ region centered on M92 with CFH12K in order to study the dynamical structure and tidal tails of the cluster. We also performed a very detailed analysis of the main sequence mass function, which was impossible to do in the study using photographic plates.  M92 (NGC 6341) is a very metal poor ([Fe/H] = $-2.24$)globular cluster located at a distance from the galactic center of $R_{GC} = 9.1$ kpc, and is $Z_{GC} = 4.3$ kpc above the galactic plane (Harris 1996).  This paper is organized as follows. Sect. 2 presents the data and the reduction procedures. In Sect. 3, we show CMDs and discuss the blue straggler stars. The radial profiles of stars in different magnitude ranges are obtained in Sect. 4. The resulting LFs and MFs are presented and discussed in Sect. 5, and the photometric mass is estimated in Sect. 6. In Sect. 7, the surface density maps are discussed. Final results are summarized in Sect. 8. ",
        "conclusions": "We have carried out a wide field CCD survey of M92 and investigated the dynamical structure and stellar population of the cluster M92. The main results of our study are:  (i) We find some candidates of BS stars in the CMD. We can not find any evidence that the BS stars are more centrally concentrated than SGB stars.  (ii) The radial density profile of M92 reveals the presence of extra tidal material. From the King model fitting method, we derived a tidal radius for M92 of $r_t = 840\\arcsec$. We find that the extra-tidal density profile is shallower for the bright star when compared to the fainter stars.  (iii) We show that the slope of cluster MF increases toward the outer part of the cluster, as expected from mass segregation. The relatively steep MF slope of M92, especially in the outer part of the cluster, suggests that this cluster has not been strongly affected by tidal shocks.  (iv) From an estimation of the photometric mass of M92, we conclude that the maximum value of the cluster's MS mass is $\\sim 1.6\\times 10^5 M_\\odot$. At least 25\\% of the total cluster mass may  be contributed by the remnant population.  (v) From the surface density map, we see a slight elongation of the cluster halo perpendicular to the direction to the Galactic center. Possible explanations of the weak tidal halo of M92 are that the tidal tail may be compressed along the line of sight, or the evaporating stars may not yet have been formed a tidal stream.  Finally, we note the need for deeper CCD photometry to investigate extra-tidal stars, especially for clusters, like M92, which have weak tidal extensions."
    },
    "0212/astro-ph0212064_arXiv.txt": {
        "abstract": "We present recent observations of small hydrocarbons (C$_3$H$_2$, C$_2$H, C$_4$H) with high abundances in the Photo-Dissociation Region of the Horsehead nebula. Our results show for the first time observational indications that the small hydrocarbon distribution follows the Aromatic Infrared Bands (AIBs) emission traced by ISO-LW2 (5-8.5\\,$\\mu$m), whereas it does not coincide with the CO and isotopes large-scale distribution. The derived abundances are significantly higher than in local clouds. This enhancement might be explained by an {\\it in situ} formation assisted by the release of carbonaceous molecules from UV-irradiated aromatic particles. ",
        "introduction": "\\label{intro}  The spatial distribution of the Aromatic Infrared Bands emission features (AIBs: 3.3, 6.2, 7.7, 8.6, 11.3 and 12.7 $\\mu $m features) has been mapped at high spatial resolution by ISOCAM in many Photo-Dissociation Regions (PDRs). The identification of the AIBs carriers is still debated, however it is well established that they are complex carbon compounds heated transiently by the absorption of a single UV or visible photon. The best candidates are Polycyclic Aromatic Hydrocarbons (PAHs: e.g. L{\\'e}ger \\& Puget 1984) or Hydrogenated Amorphous Carbons (HACs: Duley \\& Williams 1988). These large-sized carbonaceous molecules ($\\sim$ 50 carbon atoms) are likely to be chemically connected to the numerous small hydrocarbon detected the ISM, including cyanopolyynes or acetylenic chains and rings.  Various theoretical works and laboratory experiments support this point: modelling of dark cloud chemistry (Herbst 1991 ; McEwan et al. 1999), calculation of the acetylene formation rate by photo-erosion of UV irradiated PAHs (Allain et al. 1996), experiments of carbon chains production during the irradiation of HACs by a UV laser pulse (Scott et al. 1997). Nevertheless we still miss a firm observational link between AIB carriers and small hydrocarbons. PDRs are the best places to look for such a link as they exhibit strong AIB emission features which have been mapped at arcsec resolution by ISO in several objects. They are moreover highly structured regions where, due to the penetration depth of UV irradiation, the aromatic emission, the H$_2$ ro-vibrational lines, and the molecular (CO) emission are spatially separated (see Hollenbach \\& Tielens 1997). ",
        "conclusions": "\\begin{figure*} \\begin{center} \\includegraphics[width=0.8\\linewidth]{./teyssierd2.eps} \\caption{\\label{models} Abundances of various species derived from an updated Le Bourlot et al.'s PDR model, using two chemistry networks (calculations courtesy of E. Roueff).} \\end{center} \\end{figure*}  \\begin{table*} \\begin{center} \\begin{tabular}{l c c| c| c c} \\hline \\hline &  Observed  &           &  \\multicolumn{3}{c}{Column Density (\\cmd)} \\\\ \\cline{4-6} Molecule &  Frequency & Position  &  Measured    &  \\multicolumn{2}{|c}{Models} \\\\ &  (GHz)      &           &              &  $A_V^{\\rm peak} = 3$          &  $A_V^{\\rm peak} = 15$  \\\\ \\hline C$_2$H   & 87.360   & AIB peak   &  1.8$\\pm0.2$$\\times$$10^{14}$ &  1.9--5$\\times$$10^{13}$     &  1--2.5$\\times$$10^{14}$      \\\\ &          & CO peak   &  1.3$\\pm0.1$$\\times$$10^{14}$ &               &               \\\\ C$_3$H$_2$ & 85.339 & AIB peak   &  1-2$\\times$$10^{13}$      &  1--3$\\times$$10^{11}$  &  0.5--1.5$\\times$$10^{12}$    \\\\ &          & CO peak   &  0.7-1.4$\\times$$10^{13}$  &                  &               \\\\ C$_4$H   & 85.634   & AIB peak   &  2$\\times$$10^{13}$       &  1.7--3$\\times$$10^{11}$ &  0.9--1.5$\\times$$10^{12}$    \\\\ &          & CO peak   &  1.3$\\times$$10^{13}$      &                  &               \\\\ C$^{18}$O &         & AIB peak   &  4$\\pm0.2$$\\times$$10^{15}$&  5.4$\\times$$10^{14}$   &  2.7$\\times$$10^{15}$ \\\\ &          & CO peak   &  5$\\pm0.2$$\\times$$10^{15}$&                  &               \\\\ \\hline \\end{tabular} \\end{center} \\caption{\\label{tab coldens} Molecular column densities measured in the PDR (use of LTE and LVG approximations), and inferred from models (Fig.~\\ref{models}) for two peak extinctions, in areas corresponding to the AIB and the CO emission peaks. Assumptions are: $n$(H$_2) = 10^4$\\,cm$^{-3}$, $T_{\\rm ex} \\sim 10$\\,K for all hydrocarbons, and $T_{\\rm kin} \\sim 35$\\,K for C$^{18}$O.} \\end{table*}   We have compared our measurements to PDR models using several chemistry networks (Fig.~\\ref{models}) and assuming representative conditions ($G_0=100$, $n_{\\rm H}=2\\times10^{4}$\\,cm$^{-3}$). We find the following results: \\begin{itemize} \\item the models reproduce well the abundance ratios of [C$_3$H$_2$]/[C$_4$H] and [C$^{18}$O]/[C$_2$H] ($\\sim 1$ and 20 resp.) \\item  C$_3$H$_2$ and C$_4$H are at least one order of magnitude more abundant than expected compared to C$_2$H \\end{itemize}  Part of the explanation possibly lies in the under-estimated H$_2$ column density. Kramer et al. (1996) reported $A_V=3$ inside the nebula (2$^{\\prime}$ resolution), translating into molecular column densities 10-100 times lower than observed. Recent continuum observations performed at 1.2\\,mm (11$^{\\prime\\prime}$ resolution) suggest visible extinctions more likely in the range $A_V=10-20$, which reconciles the observed C$^{18}$O and C$_2$H results with the model calculations. Nevertheless, a residual over-abundance (factor $\\sim 20$) remains for the more complex hydrocarbons, suggesting that the hydrocarbon production is enhanced in the illuminated front. {\\bf We propose that this formation process is related to the photo-erosion of UV-irradiated large carbonaceous compounds at the border of the PDR}. In this scenario, complex hydrocarbons such as C$_4$H would survive the intense illumination conditions encountered at the molecular cloud interfaces. This also suggests that heavier species (C4, C5) could be observed in UV-irradiated regions."
    },
    "0212/astro-ph0212314_arXiv.txt": {
        "abstract": "One of the central problems in supernova theory is the question how massive stars explode. Understanding the physical processes that drive the explosion is crucial for linking the stellar progenitors to the final remnants and for predicting observable properties like explosion energies, neutron star and black hole masses, nucleosynthetic yields, explosion anisotropies, and pulsar kicks. In this article we review different suggestions for the explosion mechanism and discuss the constraints that can or cannot be deduced from observations. The prompt hydrodynamical bounce-shock mechanism has turned out not to work for typical stellar iron cores and empirical values of the compressibility of bulk nuclear matter. Magnetohydrodynamical models on the other hand contain a number of imponderabilities and are still far behind the level of refinement that has been achieved in nonmagnetic simulations. In view of these facts the neutrino-driven mechanism must still be considered as the standard paradigm to explain the explosion of ordinary supernovae, although its viability has yet to be demonstrated convincingly. Since spherically symmetric models do not yield explosions, the hope rests on the helpful effects of convection inside the nascent neutron star, which could boost the neutrino luminosity, and convective overturn in the neutrino-heated region behind the stalled shock, which increases the efficiency of neutrino-energy transfer in this layer. Here we present the first two-dimensional simulations of these processes which have been performed with a Boltzmann solver for the neutrino transport and a state-of-the-art description of neutrino-matter interactions. Although our most complete models fail to explode, convection brings them encouragingly close to a success. An explosion could be obtained by just a minor modification of the neutrino transport, in which case the exploding model fulfills important requirements from observations. We discuss necessary improvements on the route to finally successful models. ",
        "introduction": "\\label{sec:introduction}  Supernova explosions of massive stars are powered by the gravitational binding energy that is released when the initial stellar core collapses to a compact remnant and its radius shrinks from typically a few thousand kilometers to little more than ten kilometers. For solar-metallicity progenitors with main-sequence masses of less than about 20--25$\\,$M$_{\\odot}$ the compact leftover will be neutron star. In case of more massive stars a black hole will be formed, most likely by the fallback --- on a timescale of seconds to hours --- of matter that does not become unbound in the stellar explosion. But also the direct collapse of the stellar core to a black hole on a multiple of the dynamical timescale is possible (\\cite{wooheg02,hegetal02}).  Since the collapse proceeds essentially adiabatically the total energy of the stellar core is conserved during the implosion. The gravitational energy is temporarily stored as internal energy, mainly of degenerate electrons and electron neutrinos. If rotation plays a significant role in the progenitor core, a major fraction of the potential energy may also be converted to rotational energy of the nascent neutron star (or black hole).  The disruption of the massive star in a supernova explosion now means that some fraction of the energy in these reservoirs has to be transferred from the compact central object to the outer stellar layers. What are the physical mechanisms to mediate this energy transfer to the ejecta? And on what timescale do they work? Proposals in the literature include the hydrodynamical bounce-shock, neutrinos, or magnetic fields. The former would initiate the explosion on a dynamical timescale, whereas the latter two can establish the energy transfer only on the secular timescales of neutrino diffusion or magnetic field amplification, respectively.  Unfortunately, observations have so far not been able to yield direct insight into the processes in the stellar center at the onset of the explosion. The hope is that a future Galactic supernova will change this situation by allowing the measurements of a large number of neutrinos and possibly of a gravitational wave signal in great detail. The few neutrino events discovered in connection with Supernova~1987A were a clear signal of stellar core collapse and neutron star formation, but they were not sufficient to reveal the secrets of the explosion. Up to now we have to exploit the less direct information that is provided by the supernova light, by the characteristic properties of supernovae and diffuse and compact supernova remnants, and by the nucleosynthesis of heavy elements which takes place in the immediate vicinity of the newly formed neutron star.  Section~\\ref{sec:observations} will discuss the constraints for the explosion mechanism that are associated with such observations. In Sect.~\\ref{sec:theory} we shall briefly review the different suggestions that have been brought forward to explain the explosions of massive stars and will critically evaluate our knowledge on grounds of theoretical considerations. In Sect.~\\ref{sec:neutrinoexplosions} we shall summarize the status of detailed hydrodynamical supernova simulations and their implications for our understanding of the delayed explosion mechanism by neutrino-energy deposition behind the supernova shock. In Sect.~\\ref{sec:results} we shall present the first results of a new generation of multi-dimensional supernova models which employ a Boltzmann solver for the neutrino transport and a state-of-the-art description of neutrino-matter interactions. Section~\\ref{sec:conclusions} will conclude with an appraisal of the results and an outlook on open ends. ",
        "conclusions": "\\label{sec:conclusions}   We believe that our 2D models with a Boltzmann solver for the neutrino transport and a state-of-the-art description of neutrino-matter interactions have considerably reduced the uncertainties associated with the treatment of the neutrino physics in previous multi-dimensional simulations. With the most complete implementation of the transport physics we could not obtain explosions. This result suggests that the neutrino-driven mechanism fails with the employed input physics, at least in case of the considered 15$\\,$M$_{\\odot}$ star. We do not think that the approximate treatment of general relativistic effects is likely to jeopardize this conclusion. A comparison with fully relativistic one-dimensional calculations (Liebend\\\"orfer et al., in preparation) is very encouraging. Because of the remarkable similarity of the shock trajectories of different progenitors in spherical symmetry (\\cite{liemes02}), it is probable that our negative conclusion is also valid for other pre-collapse configurations with a similar structure. Significant star-to-star variations of the progenitor properties with a non-monotonic dependence on the stellar mass (\\cite{wooheg02}), however, suggest that multi-dimensional core-collapse simulations of a larger sample of stars are needed before one can make final, more generally valid statements.  Another concern may be the omission of lateral neutrino fluxes in our multi-dimensional transport approximation. We describe the neutrino losses from the hot spots that appear at locations where the low-entropy downflows plunge into the neutrinospheric layer, by just a radial flux. In reality the neutrinos produced in this semi-transparent matter would be radiated away more isotropically, a fact which could lead to a larger reabsorption probability in the surrounding medium. The corresponding increase of the heating efficiency might not be very large, but could still be relevant because of the close proximity of our models to an explosion. The supernova problem is highly nonlinear and surprises may lurk behind every corner.  It would therefore be premature to conclude that the neutrino-driven mechanism fails and that not even postshock convection can alter this unquestioned outcome of all current spherical models. Besides studying other progenitors with multi-dimensional simulations, one should also investigate the effects of rotation during the neutrino-heating phase and long-time post-bounce evolution of a supernova. Previous simulations found weaker explosions in case of rotating models (\\cite{fryheg00}). But these calculations were performed with a much simplified treatment of the neutrinos and showed powerful, nearly prompt explosions already in the absence of rotation. It cannot be excluded that even a moderate amount of rotation might have significant consequences when added to nonrotating models which explode only with a long time delay after core bounce or even fail to do so. Naturally, three-dimensional simulations are ultimately desirable and could reveal generic differences compared to 2D, in particular in combination with rotation. The recent results of \\hbox{(\\cite{frywar02})}, which suggest a large degree of similarity between 2D and 3D results, are groundbreaking but again must be interpreted with caution. The serious approximations of the neutrino physics and the rapid development of the explosion in these models do not allow too far-reaching conclusions.  Another major uncertainty of current supernova models is associated with the incompletely known physics in the nuclear and supranuclear medium. Not only does the high-density EoS affect the core bounce, shock formation, and prompt shock propagation, it also determines the internal properties of the hot neutron star and thus its neutrino emission, both indirectly and directly. Hadronic degrees of freedom in addition to nucleons in the supranuclear matter, e.g., hyperons, pions or kaons, can reduce the stiffness of the EoS at very high densities. The possible effects on the supernova explosion are essentially unexplored. A softening of the EoS can, for example, lead to a faster contraction and heating of the nascent neutron star. This in turn also determines the density and temperature profile in the vicinity of the neutrinosphere where most of the neutrino flux is produced during the crucial phase of delayed shock revival.  Our successfully exploding 2D model, Model~s15Gio\\_2d.a, at least demonstrates that simulations which include the effects of postshock convection are rather close to an explosion. Therefore modest changes of the neutrino emission and transport (in our case the omission of the velocity-dependent terms in the neutrino-momentum equation) seem to be already sufficient to push them beyond the critical threshold. The properties of the explosion in this case are very encouraging and may support one's belief in the basic viability of the delayed explosion mechanism. At 440$\\,$ms after bounce the shock has arrived at a radius of more than 3000$\\,$km and is expanding with about 10000$\\,$km/s (Fig.~\\ref{fig:shockradii}). The explosion of this model does not become very energetic (Fig.~\\ref{fig:explosionenergy}). The energy is just $\\sim 4\\times 10^{50}\\,$erg at that time and grows only very slowly. This may not be a serious problem if one recalls the large spread of energies of observed supernovae (Supernova~1999br, for example, is estimated to have an ejecta mass of $14\\,$M$_{\\odot}$ and an explosion energy of about $6\\times 10^{50}\\,$erg; \\cite{ham02}). Moreover, the explosion energy is a very steep function of the physical conditions in models which are above the explosion threshold (\\cite{janmue95}). A significant increase of the energy may therefore not require large changes.  Since the explosion starts rather late (at $\\sim 150\\,$ms post bounce), the proto-neutron star has accreted enough matter to have attained an initial baryonic mass of 1.4$\\,$M$_{\\odot}$. Therefore our simulation does not exhibit the problem of previous successful multi-dimensional calculations which produced neutron stars with masses on the lower side of plausible values ($\\sim 1.1\\,$M$_{\\odot}$). Also another problem of published explosion models (e.g., \\cite{herben94,burhay95,janmue96,fry99}) has disappeared: The ejecta mass with $Y_e \\la 0.47$ is less than $\\sim 10^{-4}\\,$M$_{\\odot}$ in our calculation (Fig.~\\ref{fig:yemass}), thus fulfilling a constraint pointed out by \\cite{hofwoo96} for supernovae if they should not overproduce the $N=50$ (closed neutron shell) nuclei, in particular $^{88}$Sr, $^{89}$Y and $^{90}$Zr, relative to the Galactic abundances. Of course, final statements about explosion energy, ejecta composition, and the neutron star mass (which may change due to later fallback, especially when the explosion energy remains low) require to follow the evolution for a longer time."
    },
    "0212/astro-ph0212408_arXiv.txt": {
        "abstract": "{We present a long \\bsax observation of Abell 754 that reports a nonthermal excess with respect to the thermal emission at energies greater than $\\sim$ 45 keV. A VLA radio observation at 1.4 GHz definitely confirms the existence of diffuse radio emission in the central region of the cluster, previously suggested by images at 74 and 330 MHz (Kassim \\etal 2001), and reports additional features. Besides, our observation determines a steeper radio halo spectrum in the 330-1400 MHz frequency range with respect to the spectrum detected at lower frequencies, indicating the presence of a spectral cutoff. The presence of a radio halo in A754, considered the prototype of a merging cluster, reinforces the link between formation of Mpc-scale radio regions and very recent or current merger processes. The radio results combined with the hard X-ray excess detected by \\bsax give information on the origin of the electron population responsible for nonthermal phenomena in galaxy clusters. We discuss also the possibility that 26W20, a tailed radio galaxy with BL Lac characteristics located in the field of view of the PDS, could be responsible for the observed nonthermal hard X-ray emission. ",
        "introduction": "Shocks and turbulence associated with a major cluster merger event could provide the ingredients necessary to the formation of extended radio regions detected so far in a limited number of clusters, namely a diffuse magnetic field amplification and particle reacceleration (Tribble 1993; Roettiger, Burns, \\& Stone 1999; Roettiger, Stone, \\& Burns 1999; Brunetti \\etal 2001). The existence of Mpc-scale radio halos or relics combined with the relatively short radiative lifetimes suggests an in-situ electron reacceleration induced by a very recent or current merger event. This hypothesis is supported by observational evidence (Feretti 1999). The existence of large radio regions could be at the origin of the nonthermal hard X-ray (HXR) emission detected in the Coma cluster (Fusco-Femiano \\etal 1999; Rephaeli, Gruber, \\& Blanco 1999) and Abell 2256 (Fusco-Femiano \\etal 2000) both showing extended radio emission. A considerable fraction of the input energy during a merger process can be released in particle acceleration and re-emitted in various energy bands. In particular, the same accelerated electrons responsible for the diffuse radio synchrotron emission can scatter the Cosmic Microwave Background (CMB) photons to produce nonthermal inverse Compton (IC) X-ray emission. Alternative relativistic particles to primary re-accelerated electrons emitting in radio halos or relics and responsible for IC nonthermal HXR emission could be secondary electrons produced in the interactions of cosmic rays in the ICM (Dennison 1980; Blasi \\& Colafrancesco 1999).  X-ray and optical observations report a violent merger event in the galaxy cluster Abell 754 (Henry \\& Briel 1995; Zabludoff \\& Zaritsky 1995; Henriksen \\& Markevitch 1996; Bliton \\etal 1998; Markevitch \\etal 1998; De Grandi \\& Molendi 2001) and a numerical hydro/N-body model of this cluster (Roettiger, Stone, \\& Mushotzky 1998) has shown that many of its observed morphological properties can be explained by a very recent merger ($<$ 0.3 Gyr), slightly off-axis, between two clusters having a total mass-ratio less than 2.5:1. Therefore, the intracluster medium (ICM) of A754 appears to be a suitable place for the formation of radio halos or relics. As a consequence, radio and HXR observations of this cluster are relevant to establish robust evidence of the suggested link between the presence of nonthermal phenomena and merger activity in clusters of galaxies and to obtain information on the processes that can inject relativistic electrons. The cluster has been recently imaged at 74 and 330 MHz with the NRAO VLA observatory (Kassim \\etal 2001) suggesting the existence of a radio halo and at least one radio relic. Observations at higher resolution and sensitivity are required to confirm the obtained results. A754 was observed in hard X-rays with \\rxte in order to search for a nonthermal component (Valinia \\etal 1999). The fit to the PCA and HEXTE data set an upper limit of $\\sim$ 1.4$\\times$ 10$^{-12}\\erg$ in the 10-40 keV band to the nonthermal emission.  In this paper we present the results of a high sensitivity VLA radio observation at 1.4 GHz and of a long \\bsax observation, exploiting the unique capabilities of the PDS (Frontera \\etal 1997) to search for HXR emission. We discuss the new radio data and the origin of the excess detected at energies above $\\sim$ 45 keV with respect to the thermal emission.  Throughout this paper, we assume a Hubble constant of $\\h0$ and $q_0$ = 1/2, so that an angular distance of $1^{\\prime}$ corresponds to 86 kpc ($z_{A754}$ = 0.054; Bird 1994). Quoted confidence intervals are at a 90\\% level, if not otherwise specified. ",
        "conclusions": "Our VLA observation at 1.4 GHz definitely confirms diffuse radio emission in A754 composed by a radio halo and a relic, although a connection relic-halo cannot be excluded. The relic image at 1.4 GHz appears much more extended than that detected at 74 MHz by Kassim \\etal (2001). Our observation determines with accuracy the slope of the radio halo spectrum in the 0.3-1.4 GHz frequency range that results to be steeper than that obtained in the band 74-330 MHz, indicating the presence of a radio spectral cutoff. A \\bsax observation of A754 shows a nonthermal HXR excess at energies above $\\sim$ 45 keV with respect to the thermal emission, at a confidence level of $\\sim 3\\sigma$. One possible explanation is IC emission for the same relativistic electrons responsible for the diffuse radio emission. This interpretation has been invoked to explain the \\nt emission detected in Coma and A2256 both showing extended radio regions. The lack of significative variability of the \\nt flux detected in A754 is consistent with a diffuse mechanism for the observed excess. This detection, that needs to be confirmed by future observations, and the radio results seem indicate that primary re-accelerated electrons and not secondary electrons are at the origin of the observed nonthermal emission in clusters of galaxies. At present, it exists also the possibility that the radio galaxy 26W20, located in the field of view of the PDS, is responsible for the reported hard excess. The extrapolation of the \\rosat flux in the PDS energy range requires a flat index to explain the nonthermal HXR flux. However, considering the non simultaneity of the \\rosat and \\bsax observations the SED may have assumed a different shape at the moment of the \\bsax observation. To discriminate between these two interpretations a hard X-ray observation with imaging instruments is necessary."
    },
    "0212/hep-ph0212147_arXiv.txt": {
        "abstract": "{We use the recently reported KamLAND measurements on oscillations of reactor anti-neutrinos, together with the data of previously reported solar neutrino experiments, to show that: (1) the total $^8$B neutrino flux emitted by the Sun is $1.00(1\\pm 0.06)$ (1$\\sigma$) of the standard solar model (BP00) predicted flux, (2) the KamLAND measurements reduce the area of the globally allowed oscillation regions that must be explored in model fitting by six orders of magnitude in the $\\Delta m^2-\\tan^2 \\theta$ plane, (3) LMA is now the unique oscillation solution to a CL of $4.7\\sigma$, (4) maximal mixing is disfavored at $3.1\\sigma$, (5) active-sterile admixtures are constrained to $\\sin^2\\eta\\leq 0.13$ at $1\\sigma$, (6) the observed $^8$B flux that is in the form of sterile neutrinos is $0.00^{+0.09}_{-0.00}$ (1$\\sigma$), of the standard solar model (BP00) predicted flux, and (7) non-standard solar models that were invented to avoid completely solar neutrino oscillations are excluded by KamLAND plus solar data  at $7.9\\sigma$. We also refine quantitative predictions for future $^7$Be and $p-p$ solar neutrino experiments.} ",
        "introduction": "\\label{introduction}  KamLAND is the first experiment to explore with a terrestrial beam the region of neutrino oscillation parameters that is relevant for solar neutrino oscillations. This spectacular achievement~\\cite{kamlandannouncement} makes possible new studies of the physics of neutrinos ~\\cite{sterile,friedland,cptconserv,beacom,shrock,aliani-1,grimus, holanda,bandyo} and of the solar interior~\\cite{cnopaper}. We concentrate here on refining the allowed regions in neutrino oscillation space and on what can be learned about the total $^8$B solar neutrino flux, as well as the sterile neutrino component of the $^8$B neutrino flux. The procedures we use have been described in refs.~\\cite{sterile,postsno02}.  We assume throughout this paper that CPT is satisfied and that the anti-neutrino measurements by KamLAND apply directly to the neutrino section. In fact, the first KamLAND results already show that CPT is satisfied in the weak sector to a characteristic accuracy of about $10^{-20}$ GeV~\\cite{cptconserv}.  We begin in section~\\ref{sec:kamland} by determining the allowed oscillation regions using the recently released KamLAND data. Our principal results, which are in good agreement with the analysis of the KamLAND collaboration~\\cite{kamlandannouncement}, are summarized in figure~\\ref{fig:kamlandonly} and in table~\\ref{tab:bestfits}.  We also show that the numerical values of the physical quantities that we calculate in this paper are insensitive to the details of the analysis of the KamLAND data (see discussion at the end of section~\\ref{sec:summary}).    We derive in section~\\ref{sec:global} a global solution for the neutrino oscillation parameters using, together with the KamLAND measurements, all the available solar neutrino data from the chlorine~\\cite{chlorine}, gallium~\\cite{sage02,gallex,gno}, Super-Kamiokande~\\cite{superk}, and SNO~\\cite{snoccnc,snodaynight} experiments. The enormous reduction in the size of the allowed oscillation regions is evident from a comparison of the `Before' and `After' allowed regions, figure~\\ref{fig:before} and figure~\\ref{fig:after}.  We determine in section~\\ref{sec:total} the total $^8$B solar neutrino flux, as well as the uncertainties in the flux that are implied by the existing experimental data. This flux can be compared directly with the flux predicted by the standard solar model.  Could there be a large but unobservable flux of $^8$B solar neutrinos that reaches Earth in the form of sterile (right-handed) neutrinos? Prior to the announcement of the KamLAND results~\\cite{kamlandannouncement}, the answer to this question was a resounding ``Yes\"\\cite{sterile,bargersterile}. At $3\\sigma$, the sterile component of the $^8$B flux could be, prior to KamLAND, as large as 1.1 times the BP00 predicted flux, which amounts to 50\\% of the total flux for this extreme case. However, the KamLAND measurement allows us to place a strong upper limit on the flux of sterile $^8$B neutrinos. We determine in section~\\ref{sec:sterile} the best-estimate sterile $^8$B neutrino flux and the associated uncertainties in the sterile flux.  We summarize and discuss our main conclusions in section~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  The allowed oscillation regions that are determined by including the initial KamLAND results~\\cite{kamlandannouncement} together with the previously known solar neutrino results (see figure~\\ref{fig:after}) populate 6 orders of magnitude less area in neutrino oscillation parameter space ($\\Delta m^2, \\tan^2\\theta$) than the allowed oscillation regions that are determined with just solar neutrino data (see figure~\\ref{fig:before}). The LMA solution is now the only allowed solution at $3\\sigma$. The LOW, Vacuum, and SMA solutions are disfavored at, respectively, $4.7\\sigma$, $4.8\\sigma$, and $6.1\\sigma$. All alternative explanations to the solar neutrino problem that do not also predict a deficit in KamLAND are now excluded at 3.6$\\sigma$. Combining solar neutrino data with the KamLAND measurements,  non-standard solar model explanations of the solar neutrino problem that do not include new particle physics (e.g., neutrino oscillations) are excluded at the $7.9\\sigma$ CL. Maximal mixing, $\\tan^2 \\theta = 1$, is disfavored at $3.1\\sigma$. The results in this paragraph are described in sec~\\ref{subsubsec:globalallowedregions}.  At the present level of accuracy, matter effects are not important in determining the allowed oscillation regions permitted by the KamLAND experiment. However, when the precision of the experiment is improved, matter effects must be included (see sec.~\\ref{subsec:mattereffects}). Within the currently allowed oscillation regions, the maximum change is 4\\% in the calculated event rate for the KamLAND experiment for a specific value of $\\Delta m^2$ and $\\tan^2 \\theta$.   The LMA prediction of the $\\nu-e$ scattering rate for $^7$Be solar neutrinos is precise, $0.64\\pm 0.03$ within the $1\\sigma$ (2 dof)currently allowed oscillation region (see section~\\ref{subsubsec:futureexperiments}), in units of the rate predicted by the standard solar model and assuming no flavor conversion. The estimated uncertainty in the predicted standard solar model $^7$Be solar neutrino flux, $\\pm 10$\\%~\\cite{bp00}, is much larger than the current uncertainty in the predicted scattering rate due to neutrino oscillation parameters. Thus BOREXINO and KamLAND can provide a crucial test of solar model predictions by measuring accurately the $^7$Be $\\nu-e$ scattering rate.  Table~\\ref{tab:predictions} of section~\\ref{subsubsec:futureexperiments} summarizes the predicted rates for $\\nu-e$ scattering by $p-p$ as well as $^7$Be ( and $^8$B) neutrinos. Within the $1\\sigma$ (2 dof) allowed region, the predicted $p-p$ scattering rate is $\\sim \\pm 2$\\% due to imperfectly known oscillation parameters, which is about twice the currently estimated uncertainty in the solar model prediction of the $p-p$ neutrino flux. If our current ideas regarding the accuracy of the solar model predictions are valid and if the popular neutrino oscillation framework is correct, then a measurement of the $p-p$ solar neutrino will provide a precise determination of the mixing angle $\\theta$.  It will be difficult with the existing SNO detector~\\cite{snoccnc,snodaynight} to measure the predicted SNO CC day-night asymmetry: $A_{\\rm N-D}({\\rm SNO~CC}) ~=~ 3.3^{+0.7}_{-0.6} \\% (~1\\sigma )$ (see eq.~\\ref{eq:CCdaynight}).  The event rate measured at KamLAND, (eq.~\\ref{eq:kamlandrate}) lies within the 1-sigma pre-KamLAND expectations  based on purely solar neutrino experiments (eq.~\\ref{eq:kamlandratepred}). This excellent agreement indicates that the main features of solar neutrinos are beginning to be well understood.  The most accurate method to determine the total $^8$B flux is to combine together all of the solar as well as KamLAND experimental data. Using this computationally intensive method, we find (see section~\\ref{sec:total}) that the total $^8$B solar neutrino flux is $1.00\\left(1\\pm 0.06\\right)$ in units of the standard solar model predicted flux~\\cite{bp00}.  The sterile component of the $^8$B neutrino flux can be determined similarly by using all of the available solar and KamLAND data. The allowed range of the active-sterile admixture is $\\sin^2\\eta\\leq 0.13$. Correspondingly, our best estimate for the sterile component of the $^8$B solar neutrino flux is $0.00^{+0.09}_{-0.00}$ in units of the standard solar model predicted flux (see section~\\ref{sec:sterile}).     Our results for the total $^8$B solar neutrino flux, and for the sterile component of this flux, are summarized in table~\\ref{tab:bestfits}. The present determination of the total $^8$B solar neutrino flux reduces the $1\\sigma$ experimental uncertainty in this quantity by a factor of two with respect to the previous most accurate determination~\\cite{sterile}.  The experimental bounds on the total $^8$B neutrino flux are now more than a factor of two smaller that the uncertainties in the predicted standard solar model flux~\\cite{bp00}. This is the first ocassion on which the experimental errors on a solar neutrino flux are smaller than the solar model uncertainties. The uncertainty in the sterile neutrino component of the $^8$B neutrino flux represents a factor of two improvement over the previous most stringent bound.   We conclude with an important methodological question. Do the results given in this paper depend in a significant way upon details of the analysis of the KamLAND data? This is a natural question to ask since the KamLAND results represent a crucial but low statistics addition to the solar neutrino data.  Fortunately, different analysis strategies lead to essentially the same quantitative inferences regarding solar neutrino properties and predictions.  We have studied the robustness of our quantitative conclusions to different plausible variations in the statistical analysis of the KamLAND data. Most of these variations made no discernible difference in the final answers that were obtained from the global solar plus KamLAND analysis. For quantities such as the total $^8$B neutrino flux, the sterile component of the $^8$B neutrino flux, and the predictions for observables in future experiment, there were no significant differences.  For specificity, we report here the results of a $\\chi^2$ analysis of the KamLAND data which led to the largest changes of any of the alternative analysis strategies we adopted. In this $\\chi^2$ analysis, we used the first 5 energy bins above 2.6 MeV visible energy as presented by the KamLAND collaboration but collected all the events above the fifth bin ($E > 4.725$ MeV) into a single bin (or into two bins) in order to have enough events in each bin to assume Gaussian errors. We found that the allowed regions of the KamLAND-only analysis are slightly modified, but in ways that would be noticed only by the trained and expert eye (not to the casual reader). However, the effect of using a standard $\\chi^2$ with Gaussian errors is even less important for determining the global, solar plus KamLAND, allowed regions. In particular, this alternative analysis changes $f_{\\rm B,\\, total}$ by $\\pm 0.01$ ($1\\sigma$) and $f_{\\rm B,\\, sterile}$ by $\\pm 0.005$ ($1\\sigma$).  The maximum change in the range of predicted solar neutrino observables in table~\\ref{tab:predictions} occurs for the $R(^8{\\rm B})$ and amounts to $2$\\% at 1$\\sigma$ in the largest allowed value. For all other observables, the changes are less than $2$\\% in the $3\\sigma$ range.  We are grateful E. Akmedov and M. Smy for helpful comments and to G. Feldman for asking the important question of how large are the matter effects in the KamLAND experiment. JNB and CPG acknowledge support from NSF grant No. PHY0070928. MCG-G is supported by the European Union Marie-Curie fellowship HPMF-CT-2000-00516. This work was also supported by the Spanish DGICYT under grants PB98-0693 and by Generalitat Valenciana CTIDIB/2002/24.    \\clearpage"
    },
    "0212/astro-ph0212158_arXiv.txt": {
        "abstract": "This Letter investigates the amount of dust in the intergalactic medium (IGM). The dust photoelectric heating can be the most efficient heating mechanism in the IGM where the density is very small and there are a lot of hard ultraviolet photons. Comparing the observational thermal history of IGM with a theoretical one taking into account the dust photoelectric heating, we can put an upper limit on the dust-to-gas ratio, ${\\cal D}$, in the IGM. Since the rate of the dust photoelectric heating depends on the size of dust, we find the following results: If the grain size is $\\ga 100$ \\AA, ${\\cal D}$ at $z \\sim 3$ is $\\la 1/100$ Galactic value corresponding to $\\Omega_{\\rm dust}^{\\rm IGM}\\la 10^{-5}$. On the other hand, if the grain size is as small as $\\sim 10$ \\AA, ${\\cal D}$ is $\\la 1/1000$ Galactic  value corresponding to $\\Omega_{\\rm dust}^{\\rm IGM}\\la 10^{-6}$. ",
        "introduction": "There are metals in the intergalactic medium (IGM) even in low density regions such as Lyman $\\alpha$ forest (e.g., \\citealt{cow95,tel02}). Since metal and dust are relating each other, it is sure that dust grains also exist in the IGM. However, the amount of the intergalactic (IG) dust is still quite uncertain although it seems to be not so abundant. We try to put a new constraint on the amount of the IG dust by using the IGM thermal history suggested by the recent observations of Lyman $\\alpha$ forest (e.g., \\citealt{sch00}).  After the early attempts to estimate the IG extinction \\citep{eig49,hum56}, \\cite{cra73} have obtained an upper limit on the amount of the IG dust in terms of the density parameter as $\\Omega_{\\rm dust}^{\\rm IGM}(z=0) \\la 10^{-4}$, by comparing the observed spectral energy distribution of distant (but $z\\la 0.5$) giant elliptical galaxies with an average spectrum of similar nearby galaxies (see also \\citealt{nic71,tak72}).  Observations of the redshift evolution of the QSOs spectral slope provide us with another approach to the IG dust \\citep{wri81,che91}. Although Cheng et al.\\ (1991) found no evidence for an appreciable IG reddening up to $z\\sim 2$--3, this may not be so strict constraint because it is difficult to measure the UV slope of QSOs enough precisely owing to the broad and complex iron band superposing on the continua. Recently, we obtain further constraint on the IG dust from observations of high-z supernovae (SNe) \\citep{rie98,per99}.  An upper limit of the colour excess by the IG reddening is $\\langle E_{B-V} \\rangle_{z\\sim 0.5} - \\langle E_{B-V} \\rangle_{z\\sim 0.05} \\la 0.03$ mag.  By the way, \\cite{agu99} suggests that the IG dust can be ``gray'' by a selection rule in the transfer of grains from the host galaxies to the IGM. This occurs when small ($\\la 0.1 \\micron$) grains are destroyed selectively by the thermal sputtering in the hot gas halo of the host galaxies before they reach the IGM. Then, the extinction property becomes ``gray'' since the relatively large dust survives. Importantly, such ``gray'' IG dust can avoid the detection by the IG reddening survey like above.  Even if the IG dust is really ``gray'', its evidence 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}. According to the current status, {\\it COBE} data provides us with a rough upper limit on the IG dust \\citep{loe97,fer99,agu00}. On the other hand, the submm background radiation may give a more strict constraint on the IG dust. The emission from the IG dust expected from the modern cosmic star formation history with the grain transfer mechanism of \\cite{agu99} contributes to a substantial fraction ($\\ga$ 75\\%) of the measured background radiation in the submm \\citep{agu00}. It is interesting to develop the approach of \\cite{agu00} to match a strong constraint by SCUBA whose result of the number count accounts $\\sim$ 90\\% of the submm background light (\\citealt{cal01} and references therein).  In any current status, the amount of the IG dust is still highly uncertain. Hence, it is worth trying to find a new constraint on the amount of the IG dust.  In this Letter, it is demonstrated that we can obtain such a new constraint of the IG dust by comparing the observational IGM temperature with a theoretical one if we take into account the dust photoelectric heating in the theoretical model. This is possible because the dust heating is an efficient mechanism in the IGM \\citep{nat99}. Interestingly, the effect of the dust photoelectric heating depends on not only the amount but also the size of dust. Therefore, we can put a limit on the amount of the IG dust as a function of the typical size of the IG grains. ",
        "conclusions": "\\begin{table} \\centering \\begin{minipage}{140mm} \\caption{Temperature enhancement by IG dust.} \\begin{tabular}{ccccc} \\hline $a$, ${\\cal D}/{\\cal D}_{\\rm MW}$ & $z=5$ & 4 & 3 & 2 \\\\ \\hline 0.001 \\micron, 1/1000 & 1.01 & 1.02 & 1.04 & 1.09\\\\ 0.001 \\micron, 1/100  & 1.07 & 1.21 & 1.35 & 1.88\\\\ 0.01 \\micron,  1/1000 & 1.00 & 1.00 & 1.01 & 1.04\\\\ 0.01 \\micron,  1/100  & 1.01 & 1.04 & 1.13 & 1.37\\\\ 0.1 \\micron,   1/100  & 1.00 & 1.00 & 1.04 & 1.12\\\\ \\hline \\end{tabular} \\end{minipage} \\end{table}  \\begin{table} \\centering \\begin{minipage}{140mm} \\caption{Amount of the IG dust at $z\\ga2$.} \\begin{tabular}{cccc} \\hline $a$ & ${\\cal D}/{\\cal D}_{\\rm MW}$ & $\\Omega_{\\rm dust}^{\\rm IGM}$\\\\ \\hline 0.001 \\micron & $<10^{-3}$ & $<7\\times10^{-7}$\\\\ 0.01--0.1 \\micron   & $<10^{-2}$ & $<7\\times10^{-6}$\\\\ \\hline \\end{tabular} \\end{minipage} \\end{table}  In Table 2, we summarize the obtained upper limits of the amount of the IG dust at $z\\ga2$. Let us compare these limits with the IG dust amount estimated from a possible cosmic star formation history (e.g., \\citealt{mpd98}). Once we assume a cosmic star formation history, the cosmic metal amount can be estimated. Roughly, we find $\\Omega_{\\rm metal} \\sim$ a few $\\times 10^{-5}$ at $z\\sim3$ (e.g., \\citealt{agu99}). If we postulate $\\sim$ 0.5 as the metal depletion to dust grains, which number is suitable for the Milky Way, we obtain $\\Omega_{\\rm dust}\\sim 10^{-5}$ as the total cosmic dust amount at the redshift. We show in Table 2 the upper limits in terms of the density parameter converted from the dust-to-gas ratios by $\\Omega_{\\rm dust}^{\\rm IGM}(z)=\\zeta {\\cal D}_{\\rm MW} \\Omega_{\\rm b}(z)$. For the larger grain models, we obtain an upper limit, $\\Omega_{\\rm dust}^{\\rm IGM}\\sim 7\\times10^{-6}$ which is comparable to $\\Omega_{\\rm dust}$. This means that our upper limit may not deny the possibility that most of the dust grains produced in galaxies before $z\\sim 3$ escape to the IGM as an extreme case. For very small (0.001 \\micron) grain case, we obtain $\\Omega_{\\rm dust}^{\\rm IGM}\\la 7\\times10^{-7}$. Thus, a small fraction ($\\la$ 10\\%) of such very small grains can escape from the host galaxies to the IGM because of $\\Omega_{\\rm dust}^{\\rm IGM}/\\Omega_{\\rm dust}\\la0.1$. The latter conclusion is consistent with the suggestion by  \\cite{agu99}, who proposes that the IGM is very difficult to be polluted by the very small grains.  The obtained values of the heating rate by the dust photoelectric effect are consistent with those in Nath et al.~(1999). It indicates that our result is robust against uncertainties of the adopted efficiency factor $Q$, photoelectric yield $Y$, and energy distribution function of the photoelectrons, etc., because these are somewhat different between our model and Nath et al.~(1999). If we consider only graphite grains, the photoelectric heating rate is found to be about 70\\% of that of only silicate grains, i.e., silicate grains can be more efficient heating source. The temperature in the graphite case is lower than that in the silicate case, while the decrement is $\\la$ 10\\%. These results are also consistent with Nath et al.~(1999).  We shall comment some uncertainties of our analysis. First, the effect of the spectral index of $\\alpha$ is examined. As shown in Zheng et al.~(1997), for example, a softer spectrum than $\\alpha=1$ is still compatible with observations. If $\\alpha=1.5$, the heating rates by dust at $z=3$ decrease by a factor of 0.25 ($a=$ 0.1 $\\mu$m)--0.5 ($a=$ 0.001 $\\mu$m) relatively to our case of $\\alpha=1$. If $\\alpha=2$ is used, the heating rates at $z=3$ decrease by a factor of 0.1 (0.1 $\\mu$m)--0.25 (0.001 $\\mu$m). In these calculations, $\\zeta=1/100$ and $C_{\\rm RT}=3.0$ are assumed. As the dust heating becomes less efficient, a more abundant IG dust is allowed. Then the upper limits for the IG dust amount increase by a factor of 2 (0.001 $\\mu$m)--4 (0.1 $\\mu$m) for $\\alpha=1.5$ and by a factor of 4 (0.001 $\\mu$m)--10 (0.1 $\\mu$m) for $\\alpha=2$ relative to the $\\alpha=1$ case, because the dust heating rate is proportional to the dust-to-gas ratio linearly. To summarize, a softer spectral index constrains more loosely the amount of the IG dust by a factor.  It might be possible to take very large $\\alpha$. For example, we consider a background radiation field dominated by only galaxies (e.g., $\\alpha=5$). In this case, the dust heating rates at $z=3$ become 1/100 (0.1 $\\mu$m)--1/10 (0.001 $\\mu$m) of those in $\\alpha=1$. However, the thermal histories with $\\alpha=5$ never reproduce a steep rise of the observed temperatures at the He~{\\sc ii} reionization. In addition, the transfer correction should be smaller as the spectrum is softer \\citep{abe99}. Therefore, the background spectrum is unlikely to be as soft as $\\alpha=5$ after the He~{\\sc ii} reionization. We should assume a harder index.  A caveat against our results is in the correction factor of the transfer effect. Indeed, the effects of the photoelectric heating and the transfer correction can be offset. Thus, we should check this point by solving the cosmological radiative transfer in the future work.  The current work may give a hint for the origin of the large $b$ region of the observed IGM clouds. There is a possibility that the large $b$ regions are localized. If the very small grains are located at the restricted areas, a systematic heating owing to the dust photoelectrons may explain the large-$b$ at the localized region. The principal role of the photoelectric heating in the IGM temperature also indicates that the theoretical equation of state in the IGM should be reconstructed by including the photoelectric heating effect. All of these works are being developed by the authors."
    },
    "0212/astro-ph0212092.txt": {
        "abstract": "This paper presents the results of a high-resolution imaging survey, using both ground-based and \\textit{Hubble Space Telescope} images, of a complete sample of nearby barred S0--Sa galaxies in the field, with a particular emphasis on identifying and measuring central structures within the bars: secondary bars, inner disks, nuclear rings and spirals, and off-plane dust.  A discussion of the frequency and statistical properties of the various types of inner structures has already been published.  Here, we present the data for the individual galaxies and measurements of their bars and inner structures.  We set out the methods we use to find and measure these structures, and how we discriminate between them.  In particular, we discuss some of the deficiencies of ellipse fitting of the isophotes, which by itself cannot always distinguish between bars, rings, spirals, and dust, and which can produce erroneous measurements of bar sizes and orientations. ",
        "introduction": "In the past decade, high-resolution images of the centers of disk galaxies have shown that many contain distinct, small-scale structures, such as inner bars, nuclear spirals, and nuclear rings. \\citet{buta93} present a compilation of data on nuclear rings, while \\citet{friedli96b} reviews double-barred galaxies: those with a smaller, concentric bar inside the main bar.  Infrared imaging of bars-within-bars \\citep[e.g.,][]{shaw93,shaw95,friedli96a,jung97} has shown that the structures are not composed purely of dusty gas and young stars, but are also seen in the light of the old red stars that make up most of the disk mass.  As Buta \\& Crocker point out, these features can be used to probe the galaxy dynamics; they argue that nuclear rings and spirals are linked to the inner Lindblad resonance of the large-scale bar or spiral pattern.  The prevalence of small-scale stellar features tells us that the centers of many disk galaxies are dynamically ``cool'' enough to be responsive to gravitational effects on such small scales: most of the kinetic energy is in rotation, rather than random motions.  A dynamically ``hot'' bulge, where the velocity dispersion $\\sigma$ is of the same order as the rotation speed $V$, would not form bars or stellar rings.  As \\citet{kormendy93} suggested, the bright inner regions of many galaxies may be dominated by ``cool'' disks, rather than ``hot'' bulges.  When the inner features themselves contain significant mass, their non-axisymmetric gravitational pull affects the motions of stars and gas.  For example, \\citet{shlosman89} suggested that inner bars may be important in removing angular momentum from gas, feeding it into the center where it may fuel a starburst or active nucleus.  \\citet{witold02} and \\citet{shlosman02} have presented simulations of this complex flow.  To conduct a census of central structures, we observed a complete sample of nearby disk galaxies which were already known to contain a kiloparsec-scale main bar.  Selecting S0 and Sa systems to minimize the effect of dust, we used the 3.5 m WIYN telescope, and archival optical and near-IR images from the \\textit{Hubble Space Telescope} (HST), to search for multiple bars and other central features\\footnote{The WIYN Observatory is a joint facility of the University of Wisconsin-Madison, Indiana University, Yale University, and the National Optical Astronomy Observatories.  Observations with the NASA/ESA Hubble Space Telescope obtained at the Space Telescope Science Institute, operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555.}.  Our first results for three galaxies were presented in \\citet{erwin99}; \\citet[][hereafter Paper~I]{erwin-s02} discusses the types of structures that we found, their frequencies, and their relation to other properties of the galaxies.  In this paper, we discuss our observations and the analysis techniques, and provide details and measurements for the individual galaxies.  A key advantage of our approach is that we consider multiple types of inner structures (bars, disks, rings, spirals, and off-plane dust) and carefully distinguish between them.  We also show that the common technique of fitting ellipses to isophotes, while quite useful, can often lead to erroneous measurements of bar parameters such as orientation and size. ",
        "conclusions": ""
    },
    "0212/astro-ph0212228_arXiv.txt": {
        "abstract": "We estimate the signal from large scale structure at high redshifts in redshifted $21$cm line.  We focus on $z \\simeq 3$ and the $\\Lambda$CDM cosmology.  We assume that neutral hydrogen is to be found only in galaxies, and normalise the total content to the density parameter of neutral hydrogen in damped Lyman$\\alpha$ absorption systems (DLAS). We find that the {\\it rms} fluctuations due to the large scale distribution of galaxies is very small and cannot be observed at angular scales probed by present day telescopes.  We have used the sensitivity of the Giant meter-wave Radio Telescope (GMRT) for comparison.  We find that observations stretching over $10^3$hours will be required for a $3\\sigma$ detection of such objects. ",
        "introduction": "Observations of galaxies and absorption systems at high redshifts have shown that the assembly of present day galaxies started around $6 \\geq z \\geq 1$.  At higher redshifts, the fraction of matter that had collapsed to form stars was small.  By the end of the redshift range mentioned here, the rate of star formation was starting to decline and assembly of larger structures like the clusters of galaxies is the dominant process.  In order to understand the process of galaxy formation completely, it is important to study galaxies in this redshift range in as many wavelengths as possible.  In this paper we present results for the large scale distribution of H{\\/}I at $z \\simeq 3$.  More detailed results for the entire range of redshifts will be presented elsewhere.  References for earlier theoretical work as well as related observations can be found in Bagla (1999a) and Bharadwaj (2001).  Observations of galaxies and quasars at high redshifts provide convincing proof that the inter-galactic medium is ionised out to $z \\sim 6.2$ (Pentericci et al. 2002).  Thus most of the neutral hydrogen at $z \\simeq 3$ is to be found in galaxies.  We have an estimate of the total neutral hydrogen content from observations of DLAS (Storrie-Lombardi, McMahon and Irwin , 1996).  There is no observational evidence to support the hypothesis that DLAS and Lyman break galaxies (LBG) have different properties (Fynbo et al. 2002). Thus we can safely assume that neutral hydrogen is distributed uniformly amongst galaxies at $z \\simeq 3$, and that the total amount of neutral hydrogen adds up to give us the density parameter contributed by neutral hydrogen in DLAS.  \\begin{figure} \\plotone{fig1.ps} \\caption{The expected signal (dots) from largest structures at $z\\simeq 3$ as a function of angular scale and frequency width of $1$MHz.  Estimated noise level (line) for the GMRT is shown as a function of scale for the same bandwidth and $10^3$hours of integration.  We used the sensitivity of GMRT at $327$MHz, which corresponds to a redshift slightly higher than $3$.  A $3\\sigma$ detection is possible at angular scales around $5'$.} \\end{figure} ",
        "conclusions": "We have used N-Body simulations to estimate the signal in redshifted $21$cm line from large scale structure at $z \\simeq 3$.  We find that the present day telescopes can detect extreme objects in the large scale structure at these redshifts.  In our estimation, we made use of gravity only simulations and ignored any effects coming from gas physics.  These effects play a very important role in the distribution and state of gases at small scales.  However, we are interested only in gross properties and that too at large scales and there is no reason to suspect that we will get these wrong in the method that we have used as long as we restrain our estimates to scales larger than $1$Mpc (comoving).  The weakest point in our method is that we assumed that the amount of neutral hydrogen assigned to an N-Body particle did not depend on the mass of the collapsed structure that contained this particle.  We expect larger structures ($M > 10^{13}$M$_\\odot$) to be like groups of galaxies and have less neutral hydrogen by fraction.  Also, we expect smaller structures ($M < 10^{10}$M$_\\odot$) to have less neutral fraction as these will be influenced strongly by the photo-ionising background. However, the fraction of mass in haloes at these extremes is negligible and hence we do not expect these effects to invalidate our results at these redshifts.  Such effects will be very important at lower redshifts where very massive haloes are more common, or at high redshifts where the low mass haloes contain most of the mass in collapsed structures."
    },
    "0212/astro-ph0212534_arXiv.txt": {
        "abstract": "The tight correlation between the masses of central black holes and their host spheroids in nearby galaxies and active galactic nuclei (AGN) suggests that black hole growth is closely related to their spheroid formation. Based on our previous work regarding such a joint evolutionary scheme and the consequential black hole to bulge mass correlation, we use the X-ray luminosity function of AGN and the cosmological evolution rate which are from ROSAT X-ray surveys to estimate the cosmic star formation history associated with the black hole growth. By the basic assumption that the major black hole growth occurs during the luminous AGN phase, the luminosity function of AGNs as a function of redshift traces not only the accretion history of the black holes but also the cosmic star formation history of the spheroids.  Although the space density of the especially luminous QSOs is very low, we show that the total amount of star formation associated with the massive black hole growth is almost the same as that of Lyman Break Galaxies detected by the current optical deep surveys. We thus argue that the optical deep surveys may miss about half of the net star formation in our Universe. This is probably due to in part significant dust extinction as well as the small field of view of previous optical surveys which cannot sample such rare events with relatively short time scale. However, the far infrared emission from the dust heated by star formation on-going during the black hole growth could sufficiently account for the observed SCUBA number counts, and would be the probable dominating energy source of the SCUBA population. ",
        "introduction": "Recent observations reveal a tight correlation between the masses of central black holes and their host spheroids in nearby galaxies and AGNs \\citep{Kor95, Fab97, Mag98, Lao98, Ho99, Wan99a}. A mean value of the black hole to bulge mass ratio, $\\rm M_{\\rm bh}/M_{\\rm bulge}\\,\\sim\\, 10^{-3}$ and the local black hole mass density, $\\rho_{\\rm bh}\\,\\sim\\, 3 \\,-\\, 5\\,\\times\\, 10^5\\,\\,{\\rm M_{\\odot}}/{\\rm Mpc^3}$ are now fairly agreed upon by observations using different techniques \\citep{Sal99, Kor01, Mer01, McL02}. This correlation implies a possible scenario where the black hole accretion history in the inner few parsecs and the $\\,\\sim\\,\\rm kpc$ scale star formation during spheroid formation are closely connected. Simply speaking, the amount of gas accreted to grow the central black hole is proportional to the amount of gas driving the star formation to populate the spheroid within a similar time scale of the two activities \\citep{Wan98, Wan00, Hae00, Mon00, Bur01, Pag01}.  This relation also opens an important avenue to the understanding of the cosmic star formation history (CSFH) associated with the black hole growth by accretion, complementary to the previous optical/near-infrared (NIR) or far-infrared deep surveys.  The study of the CSFH from the optical/NIR deep surveys describes only the unobscured star formation in galaxies, and may be blind to activities during the massive spheroid formation because of the significant dust extinction or the small field of view \\citep{Mad96, Ste96, Con97,Ste99}. In fact, our knowledge of the cosmic star formation history, especially at high redshift, has dramatically changed in the last few years due to the results of the far-infrared and submillimeter deep surveys. The compelling lines of evidence show that the absorption and re-radiation of light by dust in the early epoch of galaxy formation and evolution is significant \\citep{Pug96, Gui97, Sch98, Fix98, Hau98, San99, Bar99}.  Although the observations at far-infrared and submillimeter wavelengths may help to unveil the early dusty star formation epoch, it is still a very difficult job to obtain an unbiased view of the early star formation history from these luminous infrared systems due to the rather poor angular resolution of SCUBA and to the faintness of the optical counterparts \\citep{Bla99a,DeZ01}.  The main purpose of this paper is to fairly reconstruct the global star formation history associated with AGN accretion and give a proper estimation of the star formation activity during spheroid formation. We assume the X-ray luminosity $\\rm L_{\\rm x}$ of the AGN is powered by accretion onto a central massive black hole, especially during the luminous phase. Deep X-ray surveys provide a direct probe of the AGN accretion history, and the history of the black hole growth as well as the joint star formation, based on the assumption that the luminous AGNs reflect the stage of major black hole growth and spheroid formation \\citep{Boy00, Miy00}.  In this paper, we trace the AGN evolution with the soft X-ray $0.5\\,-\\, 2\\,\\,{\\rm keV}$ local luminosity function and the evolution rate by Miyaji et al. (2000), which gives an excellent fit to the ROSAT surveys of various depth, including the number counts and the soft X-ray background.  The X-ray view of the CSFH is recently discussed by several authors for normal spiral galaxies where the X-ray emission is dominated by a population of X-ray binaries, hot interstellar gas, or even low luminosity AGNs (LLAGNs) with fluxes about $10^{-17}\\,-\\,10^{-18}\\,\\,{\\rm erg\\,\\,cm^{-2}\\,\\,s^{-1}}$ at $0.5\\,-\\,2\\,\\,{\\rm keV}$ energy band \\citep{Cav00, Pta01, Hor02, Miy02}. It is not the purpose of this paper to give any constraints on the star formation in galaxies or LLAGNs. Instead, we study the intensive star formation associated with black hole growth in massive spheroids using the X-ray surveys. We adopt in the calculation: 1) a simple connection between the X-ray emission from AGNs and the black hole mass by assuming an Eddington ratio $\\epsilon$; 2) the black hole to bulge mass correlation; 3) a similar time scale for the black hole growth and the intensive star formation populating finally the spheroids based on our previous work. The present day black hole density and the SCUBA number counts are used as two important model constraints in the calculation, with the set of cosmological parameters $(\\Omega_{\\rm m}, \\Omega_{\\rm \\Lambda})=(0.3, 0.7)$ and $\\rm H_{0}\\,=\\,50\\,\\,{\\rm km/s/Mpc}$. ",
        "conclusions": "We have begun to probe the cosmic star formation history associated with AGN accretion by X-ray deep surveys based on the tight correlation of $\\rm M_{\\rm bh}\\,-\\,M_{\\rm B, bulge}$ in early type galaxies and nearby AGNs. This approach is parallel and complementary to the current study from optical and infrared observations of the star formation activities at $\\rm z\\,>\\,1$ \\citep{Mad98, Pet98, Hug98, Bar98, Ell98}.  The cosmic star formation history associated with AGN accretion derived from our calculation is approximately comparable to that of the normal galaxies in Madau plot after a reasonable dust correction, although they are much rarer objects compared with optically selected galaxies (The co-moving number density of luminous AGNs is as small as $10^{-7}\\,\\,{\\rm Mpc^{-3}}$, see Fig. \\ref{fig2}). In this case, we might say that about half of the star formation (if not more) in our Universe is closely connected with AGN accretion, and we might severely overlook the intensive star formation during the epoch of the spheroidal formation by optical/NIR surveys due to the significant dust extinction and the small sample volume of these surveys.  We found from Fig. \\ref{fig2} that the peak of the intensive star formation representing the spheroid formation is at $\\rm z\\,\\sim\\, 2$, not necessarily much beyond this epoch even if we take into account a reasonable fraction of type 2 QSOs at high redshift. According to various observation, we divide the soft X-ray luminosity function into two luminosity regions, where the abundance ratio of the type 2 to type 1 Seyferts is set equal to 4, and the ratio of the type 2 to type 1 QSOs is simplified as a power law function of $\\alpha\\,(1+z)^{p}$. Actually an upper limit of the abundance ratio of the type 2 to type 1 QSOs is about $1\\,\\sim \\, 2$ from our calculation with the constraints of the local black hole mass density and the results of the submillimeter deep surveys. However, current model constraints are not sufficient to reject a much higher abundance ratio for low luminosity type 2 to type 1 AGNs ($>4$). The reason may be that the bright SCUBA counts are dominated by the luminous AGNs and the ratio of the low luminosity AGNs is not critical. Future results of the Chandra and XMM/Newton deep surveys would give more information on the abundance of the obscured objects.  The energy budget of the submillimeter sources has been discussed by several authors, which show that the AGN powered far-infrared emission in the obscured objects could account for only a certain fraction of the SCUBA number counts \\citep{Alm99, Ris02}. The SCUBA/X-ray anti-correlation of the Chandra deep surveys gives a clear detection of AGN activities in those samples only $\\,\\sim\\, 10\\%$ \\citep{Bau00, Fab00, Bar01b, Hor01}. Nevertheless, Fig. \\ref{fig3} shows that the far-infrared emission from the dust heated by the intensive star formation during the black hole growth could sufficiently interpret the number counts of the submillimeter deep surveys, where the SCUBA number counts are dominated by the contribution from the star formation in the host galaxies of the luminous type 1 or type 2 QSOs, and the the faint parts ($\\rm S_{850}< 1\\,mJy$) may be from those low luminosity X-ray sources with small bulges. In this case, our calculation may indicate that the star formation activity might dominate the energy power (at least comparable to the AGNs) in the far-infrared emission in these SCUBA sources, consistent with the multiwavelength observations of the submillimeter selected galaxies. They suggest that even when an AGN is present in a SCUBA source, it rarely dominates the engergy budget of the galaxy \\citep{Fra98, Ale02, Ivi02, Sma02}."
    },
    "0212/astro-ph0212045_arXiv.txt": {
        "abstract": "The supranova model for $\\gamma$-ray bursts (GRBs) is becoming increasingly more popular. In this scenario the GRB occurs weeks to years after a supernova explosion, and is located inside a pulsar wind bubble (PWB). Protons accelerated in the internal shocks that emit the GRB may interact with the external PWB photons producing pions which decay into $\\sim 10^{16}\\;$eV neutrinos. A km$^2$ neutrino detector would observe several events per year correlated with the GRBs. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/astro-ph0212273_arXiv.txt": {
        "abstract": "\\citet{MREG97} accurately determined the position of \\sgra\\ on an infrared image, by aligning the infrared image with positions measured for SiO masers, associated with infrared-bright evolved stars, at radio wavelengths. We now report greatly improved radio positions and, for the first time, proper motions of many stellar SiO masers at the Galactic Center. These positions and motions, coupled with better infrared imaging, allow a much improved location of \\sgra\\ on infrared images. With current data, infrared stars can be placed in a reference frame tied to \\sgra, to an accuracy of $\\approx10$~mas in position and $\\approx1$~\\masy\\ in motion. The position of \\sgra\\ is, within uncertainties, consistent with stellar accelerations and the measured orbital focus of the star S2.  The star S2 was observed within 16~mas ($\\approx130$~AU in projection) of Sgr~A* on 2002 May 02. Finally, we find that the central stellar cluster moves with \\sgra\\ to within $\\approx70$~\\kms. ",
        "introduction": "Sagittarius A* (\\sgra) is a compact radio source projected toward the center of the Galaxy \\citep{BB74}.  The apparent proper motion of \\sgra, measured against extragalactic radio sources, is consistent with it being at the dynamical center of the Galaxy \\citep{R99,BS99}. A dense cluster of stars surrounds \\sgra, and the {\\it relative} proper motions of these stars increases dramatically with decreasing projected distance from \\sgra\\ \\citep{EG96,G98}. These stellar motions exceed $1000$\\kms\\ at a projected distance of $0.015$~pc from \\sgra. Using the radio--infrared alignment presented in this paper, one star, S2, was found to travel on an elliptical orbit with Sgr A* at its focus \\citep{S02}. This star was moving at $\\gax5000$~\\kms\\ at a distance of 124~AU at pericenter passage, implying an enclosed mass of $\\approx 4 \\times 10^6$~\\msun. The very small peculiar motion of $<20$~\\kms\\ for \\sgra\\ is in stark contrast to the high stellar motions and indicates that \\sgra\\ must contain at least $10^3$ to $10^4$~\\msun\\ \\citep{R99,CHL02}; this mass is directly associated with the compact radio source which is $\\lax1$~AU in size \\citep{BB98,K98}. Together, all of these findings provide compelling evidence that \\sgra\\ is a super-massive black hole at the dynamical center of the Galaxy.  While \\sgra\\ is a very bright radio source, it is a very dim infrared source.  Locating \\sgra\\ on infrared images is complicated by the high projected density of stars at the Galactic Center and by the absence of any obvious infrared counterpart for \\sgra. Fortunately, the Galactic Center stellar cluster contains red giant stars that are both strong radio sources (from circumstellar SiO maser emission) and bright infrared sources.  Because these stars are visible at both radio and infrared wavelengths, one can use their radio positions, measured with respect to \\sgra, to determine accurate ``plate'' scales and rotations for infrared images and transfer the position of \\sgra\\ to these infrared images. In \\citet{MREG97} (hereafter Paper~I), we developed this technique and determined the position of \\sgra\\ on a diffraction-limited 2.2~\\micron\\ image of the Galactic Center region made with the MPE SHARP camera on the ESO 3.5-m New Technology Telescope in Chile. We located \\sgra\\ with a 95\\% (\\twosig) confidence of 60~mas and found no infrared emission above a level of 9~mJy (de-reddened) toward its position.  This upper limit strongly constrains models for the emission from the immediate vicinity of a super-massive black hole.  Locating \\sgra\\ on infrared images is important for reasons other than determining its emission. Recently it has become possible to observe gravitational acceleration directly, via orbital curvature in the motions of stars \\citep{G00,E02}.  Uncertainty in the position of the dynamical center, presumably \\sgra, limits estimates of the enclosed mass determined from observed accelerations. Approximately two-thirds of a complete orbit has now been observed for one star, S2 {\\citep{S02}, and calculating orbital parameters is enhanced by precise knowledge of the dynamical center. Ultimately, verification that the dynamical center coincides with the position of \\sgra\\ would link the radiative source, \\sgra, directly with the gravitational source.  In the five years since Paper I was published, there have been significant advances in both the radio astrometry and infrared imaging of the Galactic Center region, and it was time to update the infrared positioning of \\sgra. This paper reports VLA and VLBA observations with greatly improved positions and, for the first time, proper motions of stars with circumstellar SiO masers in the Galactic Center region. In addition, deep (K$_s=18$~mag), large-field images of the Galactic Center region at infrared wavelengths were obtained with new instrumentation on one of the 8.2-m VLT telescopes.  Combining the new radio and infrared observations has yielded a significant improvement in the location of \\sgra\\ in the infrared, helped verify that the radio source \\sgra\\ is at the focus of a stellar orbit, and placed the proper motions of stars in the central cluster in a reference frame tied to \\sgra. ",
        "conclusions": ""
    },
    "0212/astro-ph0212103_arXiv.txt": {
        "abstract": "The conjecture that ultra high energy cosmic rays (UHECRs) are actually strangelets is discussed. Besides the reason that strangelets can do as cosmic rays beyond the GZK-cutoff, another argument to support the conjecture is addressed in this letter via the study of formation of TeV-scale microscopic black holes when UHECRs bombarding bare strange stars. It is proposed that the exotic quark surface of a bare strange star could be an effective astro-laboratory in the investigations of the extra dimensions and of the detection of ultra-high energy neutrino fluxes. The flux of neutrinos (and other point-like particles) with energy $>2.3\\times 10^{20}$ eV could be expected to be smaller than $10^{-26}$ cm$^{-2}$ s$^{-1}$ if there are two extra spatial dimensions.  \\vspace{0.4cm} % \\noindent % {\\em PACS} numbers: 04.70.Dy, 12.38.Mh, 13.85.T, 97.60.Jd ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/astro-ph0212453_arXiv.txt": {
        "abstract": "The characteristic photon energy for Gamma Ray Bursts, $E_{peak}$, has a remarkably narrow distribution for bursts of similar peak flux, with values between 150 and 600 keV for most faint bursts.  This result is surprising within the framework of internal shock models, since spectral shifts associated with the jet's blue shift (by a Lorentz factor of $\\Gamma$) and the cosmological red shift (by a factor of $1+z$) should cause substantial smearing in the distribution of the spectral peak in the jet's co-moving frame, $E_{rest}$.  For the general case where the luminosity ($L$) varies as $\\Gamma^N$ and $E_{rest}$ varies as $\\Gamma^M$, then the observed $E_{peak}$ will vary as $L^{(M+1)/N}(1+z)^{-1}$.  For two independent set of 20 and 84 bursts, $E_{peak}(1+z)$ varies as a power law of the luminosity with an index of $(M+1)/N=0.36 \\pm 0.03$.  With this measured value, the above functional dependence of $E_{peak}$ on $L$ and $z$ results in $E_{peak}$ being roughly constant for bursts of similar peak flux, $P_{256}$.  Thus, the kinematic smearing will be small, hence allowing the $E_{peak}$ distribution to be narrow.  This model also predicts that bright bursts will have high $E_{peak}$ values because they all have some combination of high luminosity (and hence a large blue shift $\\Gamma$) and a nearby distance (and hence a small cosmological red shift).  Quantitatively, $E_{peak}$ should vary roughly as $P_{256}^{~~~0.36}$, and this model prediction is strikingly confirmed with BATSE data by Mallozzi et al.  A prediction of this model is that GRBs at very high red shift ($z \\sim 10$) should all appear with $E_{peak}$ at $\\approx 200~keV$.  A further prediction of this model is that normal bursts with $P_{256}$ below the BATSE trigger threshold will appear as x-ray flashes with $E_{peak} \\approx 70~ keV$; just as is reported by Kippen et al. and Heise et al. ",
        "introduction": "Gamma-Ray Bursts (GRBs) are indeed 'gamma-ray' bursts since their spectrum is dominated by gamma radiation.  $E_{peak}$ is the characteristic photon energy measured from a burst spectrum by fitting a smoothly broken power law (Band et al. 1993).  \\citet{mal95} have measured $E_{peak}$ values for 399 BATSE bursts.  They find that a histogram of $E_{peak}$ is roughly log-normal in shape, with average values of $\\sim 200~keV$.  They also find that the width of this distribution is surprisingly narrow, with a dispersion corresponding to a factor of $\\sim 2$ about the average. \\citet{bra98} have proven that the narrowness of this distribution is not an artifact of BATSE detector properties or of the analysis procedures, and in particular that the sharp falloff in frequency down to at least 50 keV is certain.  Further, deep searches with the Solar Maximum Mission Gamma Ray Spectrometer have proven that the population of high-$E_{peak}$ bursts is nearly zero \\citep{has98}.  This is not to say that 'GeV bursts' and 'keV bursts' do not exist. For example, GRB950425 is certainly a GeV burst \\citep{sch98}.  Also, \\citet{kip01} and \\citet{hei01} report x-ray transients that have similar light curves and durations as classical GRBs.  The existence of these x-ray transients might be interpreted as implying that the $E_{peak}$ distribution has a broad tail to low energies, despite the iron-clad data from BATSE.  The resolution lies in the realization that the x-ray transients all have peak fluxes (in the 50-300 keV band) that are far below the normal BATSE trigger threshold.  The result from \\citet{mal95} is that the $E_{peak}$ distribution is narrow {\\it for bursts in a small range of peak flux}.  Bursts with peak fluxes far below the BATSE threshold do not (and should not) enter into the histograms constructed from BATSE triggers.  If an $E_{peak}$ distribution is constructed for very low peak flux events (including the x-ray transients), then a simple extrapolation of the average $E_{peak}$ versus peak flux relation of Mallozzi suggests that such events will appear as ordinary bursts except with $E_{peak} \\sim 70$ keV, just as observed.  So apparently there is something about GRBs (the $E_{peak}$ value) that is fairly constant over all bursts.  This poses a double challenge to theory, since there must be some physical mechanism which acts as a thermostat and there must be some mechanism for balancing the kinematic blue shifts (due to the jet's velocity towards Earth) and red shifts (due to the cosmological expansion).  Based on the factor of $\\sim 1000$ variation in burst luminosity, the Lorentz factor of the jet likely varies by a large factor, while the cosmological red shift varies by over a factor of 3; so we might expect the kinematic effects to produce a broad $E_{peak}$ distribution even if all bursts had the same rest frame temperature.  \\citet{mal95} also found that the average $E_{peak}$ value varied systematically with the apparent brightness of the burst.  The peak flux used here, $P_{256}$, is measured over a 256 ms time scale within a limited energy band from 50-300 keV \\citep{fis94}.  The bright BATSE bursts have $\\langle E_{peak} \\rangle$ around 339 keV while the faintest BATSE bursts have $\\langle E_{peak} \\rangle$ around 175 keV.  The trivial explanation of this effect as arising only from cosmological red shifts is unreasonable since burst peak fluxes vary more due to the wide spread in luminosity than due to distances.  So a second mystery is why $\\langle E_{peak} \\rangle$ varies with $P_{256}$ as observed. ",
        "conclusions": ""
    },
    "0212/astro-ph0212179_arXiv.txt": {
        "abstract": "Recent results suggest that $z\\simeq 6$ marks the $end$ of the reionization era. A large sample of objects at $z\\simeq 6$, therefore, will be of enormous importance, as it will enable us to observationally determine the exact epoch of the reionization and the sources that are responsible for it. With the HST {\\it Advanced Camera for Surveys} (ACS) coming on line, we now have an unique opportunity to discover a significant number of objects at $z\\simeq 6$. The pure parallel mode implemented for the {\\it Wide Field Camera} (WFC) has greatly enhanced this ability. We present our preliminary analysis of a deep ACS/WFC parallel field at $|b|=74.4^o$. We find 30 plausible $z\\simeq 6$ candidates, all of which have $S/N>7$ in the F850LP-band. The major source of contamination could be faint cool Galactic dwarfs, and we estimated that they would contribute at most 4 objects to our candidate list. We derived the cumulative number density of galaxies at $6.0\\leq z\\leq 6.5$ as 2.3 arcmin$^{-2}$ to a limit of 28.0 mag in the F850LP-band, which is slightly higher than our prediction. If this is not due to an underestimated contamination rate, it could possibly imply that the faint-end slope of the $z\\simeq 6$ luminosity function is steeper than $\\alpha=-1.6$. At the very least, our result suggests that galaxies with $L<L^*$ do exist in significant number at $z\\simeq 6$, and that they could be the major sources that contributed the reionizing photons. ",
        "introduction": "In the last couple of years, great progress has been made in answering the grand cosmological question of when reionization of the universe began. It is now generally believed that we have seen {\\it the end} of the reionization at around $z=6$. Fan \\etal (2002) provided a detailed analysis of the spectra of the SDSS $z\\sim 6$ quasars, and argued that the epoch of reionization could not be at a redshift much higher than 6. This is also consistent with the most recent numerical simulations of cosmological reionization (e.g., Cen \\& McDonald 2002), where it is suggested that the IGM is likely neutral at $z > 6.5$. On the other hand, Hu \\etal (2002) reported their discovery of a $z=6.56$ \\Lya emitter and therefore suggested that the reionization lies beyond $z=6.6$. Nevertheless, the detectability of \\Lya emission line $prior$ to the reionization was investigated by Haiman (2002), who pointed out that such a line, if broad enough, could still be observable, and thus the \\Lya emitter of Hu \\etal (2002) might actually lie $beyond$ the epoch of the reionization. Very recently, Cen (2002) proposed a scenario that the Universe was actually reionized twice, which could also explain the seemingly-paradoxical detection of this $z=6.56$ \\Lya emitter.  The kinds of sources responsible for the reionization is one other question that remains open. While the strong UV spectrum of a quasar is capable of ionizing surrounding IGM out to Mpc scales, the paucity of quasars makes them unlikely the major contributors to the reionizing background. Based on the stringent constraint they obtained for the bright-end slope of quasar luminosity function at $z\\sim 6$, Fan \\etal (2002) pointed out that luminous quasars could not provide sufficient photons to keep the universe ionized at this redshift. However, there is still room for lower luminosity AGNs ($M_B > -23$ mag) being important reionizing sources. Alternatively, the reionizing photons could also come from star-forming galaxies, provided that a sufficient amount of Lyman continuum emission could escape from the galaxies. While the escape fraction ($f_{esc}$) of such emission is found to be extremely small in the local universe (Leitherer \\etal 1995; Hurwitz \\etal 1997), Steidel \\etal (2001) derived $f_{esc}\\simeq$ 65\\% from the composite spectrum of a sample of 29 $z\\simeq 3.4$ Lyman-break galaxies. Although such a high $f_{esc}$ value is still under debate (see Giallongo \\etal 2002), star-forming galaxies cannot be ruled out as important contributors to the reionizing background. An extreme scenario has been proposed by Ricotti (2002), who suggested that globular clusters, which can be formed outside of galaxy disks or bulges and thus have $f_{esc}$ close to unity, were responsible for the reionization.  Obviously, we will have to first acquire a large sample of $z\\simeq 6$ objects before we can settle all these issues. The {\\it Wide Field Camera} (WFC) of the {\\it Advanced Camera for Surveys} (ACS) on-board the HST provides an unique opportunity to discover a large number of $z\\simeq 6$ objects and probe the faint end of the luminosity function (LF). To maximize its scientific return, a default pure parallel observation mode has been implemented for the WFC (Sparks \\etal 2001). This mode utilizes four filters, namely F775W (SDSS-$i$), F850LP (SDSS-$z$), F475W (SDSS-$g$), and F625W (SDSS-$r$), in order of preference. Whenever there is at least one orbit of time available for parallel observation, this mode always takes images in the F775W and F850LP bands. Given the high throughput and the wide field-of-view (FOV) of the WFC, such a strategy makes its parallel data extremely useful for selecting objects at $z\\simeq 6$ by using the {\\it drop-out} technique.  In this letter, we present our preliminary analysis of a field that is the deepest ACS/WFC parallel observation to date. We describe the data reduction in \\S 2, followed by the $z\\simeq 6$ candidate selection in \\S 3. The comparison between the result of this study and the prediction of Yan \\etal (2002) is made in \\S 4. We conclude with a brief summary in \\S 5. Throughout the paper, we use the AB magnitude system, and adopt a cosmological model with $(\\Omega_M,\\Omega_\\Lambda, H_0)=(0.3,0.7,65)$. ",
        "conclusions": "The ACS/WFC provides an unprecedented opportunity to probe the LF of $z\\simeq 6$ galaxies at its faint end. The two red filters that it has, the F775W and the F850LP, are ideal for selecting such objects by identifying the Lyman-break signature at $6.0\\leq z\\leq 6.5$. We present our preliminary analysis on a deep ACS/WFC parallel field. We discovered a total of 30 plausible candidates to a limit of 28.3 mag in the F850LP-band, 27 of which are to a limit of 28.0 mag. All these candidates were detected at $S/N\\geq 7.2$ in the F850LP-band.  The contamination to our sample due to either emission-line galaxies or elliptical galaxies is likely negligible. The major contaminator would be cool Galactic dwarfs, whose effect is hard to quantify as their spatial distribution is not well known. To the best of our knowledge, we estimate that they could contribute at most 4 objects to the candidate list.  We derive a galaxy cumulative number density at $6.0\\leq z\\leq 6.5$ as 2.3 arcmin$^{-2}$ to a limit of 28.0 mag, which is slightly higher than our earlier prediction (Yan \\etal 2002). While this higher observed value could be an illusion due to an underestimated contamination of cool Galactic dwarfs, it could also be real. If the later is the case, one possible explanation is that the faint-end slope of the actual LF at $z\\simeq 6$ may be as steep as $\\alpha=-2$. Given the faintness of these candidates, spectroscopic identification is extremely difficult with existing instruments. However, deep IR imaging data now feasible with the revived NICMOS will enable us to distinguish these two possibilities. At the very least, the number of plausible candidates strongly suggests that galaxies with $L<L^*$ do exist in large number at $z\\simeq 6$, and that they are important contributors to the reionizing background. Future IR surveys with the JWST will allow us to map the LF beyond $z\\simeq 6$ and to even fainter fluxes."
    },
    "0212/astro-ph0212386_arXiv.txt": {
        "abstract": "{We present optical spectra and redshift measurements for 178 flat-spectrum objects from the Parkes quarter-Jansky flat-spectrum sample\\thanks{Table A1 is also available in electronic form at the CDS via anonymous ftp to {\\tt cdsarc.u-strasbg.fr} (130.79.128.5) or via {\\tt http://cdsweb.u-strasbg.fr/cgi-bin/qcat?J/A+A/}}. These spectra were obtained in order to compile a complete sample of quasars for use in a study of quasar evolution.  We present a composite optical spectrum made from the subset of 109 quasars that have flux densities in the range $\\rm 0.25Jy < S_{2.7GHz}< 0.5Jy$, and make a comparison with a composite for radio-quiet QSOs from the Large Bright Quasar Survey.  Our large sample of radio-loud quasars allows us to strengthen previous reports that the Ly$\\alpha$ and CIV emission lines have larger equivalent width in radio-loud quasars than radio-quiet QSOs to greater than the 3$\\sigma$ level.  However we see no significant difference in the equivalent widths of CIII] or MgII.  We also show that the flux decrements across the Lyman-$\\alpha$ line ($D_A$) measured from these spectra show the same trend with redshift as for optically selected QSOs. \\begin{keywords} quasars:general radio  \\end{keywords} } ",
        "introduction": "This is the second paper in a series of three describing the results of a program to search for high-redshift quasars and study the evolution of the flat-spectrum radio-loud quasar population.  The first paper (Paper~1, Jackson et al. 2002) set out the sample, discussing its selection and the identification and re-confirmation programme to determine the optical counterparts to the radio sources. This paper, Paper~2, presents new spectroscopic observations and redshift determinations.  Paper~3 (Wall et al. 2002) considers the quasar space distribution, the luminosity function and its epoch dependence.  Because the sample is radio selected, the completeness of the sample is unaffected by obscuration from dust. An earlier version of this sample was used by Shaver et al. (1996b) to show that the space density of quasars declines at redshifts greater than~3.  In addition to the primary goal of this work (the study of quasar evolution), the sample can be used for several other statistical studies. It has provided the basis for a study of damped Lyman-$\\alpha$ absorption systems along the lines of sight to the quasars (Ellison et al. 2001, 2002). In this paper we also consider some properties of the quasar spectra themselves, in particular the equivalent width distributions of the emission lines, and the flux decrement across the Lyman-$\\alpha$ line at high redshift. ",
        "conclusions": "We have presented spectra and redshift measurements for a large sample of radio-selected quasars which, when combined with data from the literature, forms a complete sample of radio-loud quasars. We have studied statistical properties of the new spectra, in particular the rest-frame equivalent width distribution of the rest-frame UV spectral lines and the Lyman-$\\alpha$ decrement at high redshift. We find that the spectra we have obtained show significantly stronger Ly$\\alpha$ and CIV emission lines than radio-quiet quasars from the LBQS, which has a similar optical magnitude distribution to our sample. We find no major difference in the Lyman-$\\alpha$ decrement measured from our spectra compared to those of optically-selected QSOs at similar redshifts."
    },
    "0212/astro-ph0212501_arXiv.txt": {
        "abstract": "{ We report on deep BVR-imaging of the field of the nearby millisecond pulsar PSR~J0030+0451 obtained with the ESO/VLT/FORS2. We do not detect any optical counterpart down to $B \\gtrsim 27.3$, $V \\gtrsim 27.0$ and $R \\gtrsim 27.0$ in the immediate vicinity of the radio pulsar position. The closest detected sources are offset by $\\gtrsim 3\\arcsec$, and they are excluded as counterpart candidates by our astrometry. Using our upper limits in the optical, and including recent {\\sl XMM-Newton\\/} X-ray data we show that any nonthermal power-law spectral component of neutron star magnetospheric origin, as suggested by the interpretation of X-ray data, must be suppressed by at least a factor of $\\sim 500$ in the optical range. This either rules out the nonthermal interpretation or suggests a dramatic spectral break in the  $0.003-0.1$~keV range of the power-law spectrum. Such a situation  has never been observed in the optical/X-ray spectral region of ordinary pulsars, and the origin of such a break is unclear. An alternative interpretation with a purely thermal X-ray spectrum  is consistent with our optical upper limits. In this case the X-ray emission is dominated by  hot polar caps of the pulsar. ",
        "introduction": "Millisecond pulsars (hereafter MSPs) differ from ordinary radio pulsars by much shorter spin periods $P$, smaller period derivatives $\\dot{P}$, higher dynamical ages $\\tau$, weaker magnetic fields $B$, and evolution histories (see, e.g., recent review by \\cite{Lor01}). Contrary to ordinary pulsars, only 9 of 56 MSPs currently known in the Galactic disk and 25 of 52 MSPs found in globular clusters are isolated objects (\\cite{Lor01}; \\cite{Lor02}); \\cite{Pos01}). It is believed that the fast rotation of these neutron stars (NSs) was gained in the past by angular momentum transfer during mass accretion from a companion star (\\cite{BH91}). This was supported by the discoveries of three accretion-powered X-ray MSPs in low-mass X-ray binaries (e.g., SAX J1808.4-3658, see \\cite{WvdK98}).  \\begin{table*}[t] \\caption{Parameters of \\psr\\ (Lommen et al.\\ \\cite{L00} and \\cite{Beck00}, unless specified otherwise).} \\label{t:param} \\begin{tabular}{ccccccccccc} \\hline\\hline \\multicolumn{6}{c}{Observed}&&\\multicolumn{4}{c}{Derived} \\widerul\\\\ \\cline{1-6}\\cline{8-11} $P$ & $\\dot P$ & $\\alpha_{2000}$, $\\delta_{2000}^{a,b}$ & $\\mu_\\alpha$, $\\mu_\\delta^b$ & $l$, $b^d$ & $D\\!M^e$ && $\\tau$ & $B$ & $\\dot E$ & $d^f$ \\wideru \\\\ ms & $10^{-20}$ && mas yr$^{-1}$ && cm$^{-3}$ pc && Gyr & G  & erg $s^{-1}$ & pc \\widerul \\\\ \\hline 4.865 & $1.0\\pm0.2$ & \\twolines{$00^h30^m27$\\fs432(9)$^c$}{04\\degr51\\arcmin39\\farcs67(2)} & \\twolines{$4.1\\pm3.4$}{$-24.6\\pm8.0$} & \\twolines{113\\fdg1}{$-$57\\fdg6} & $4.3328(2)$ && 7.7 & $2.2 \\times 10^{8}$ & $3.4 \\times 10^{33}$ & 230 \\widerul\\\\ \\hline \\end{tabular} $^a$~Coordinates are at the epoch of the VLT observations, MJD 52134 (Aug 13, 2001) \\ \\ \\ \\ \\ $^d$~Galactic coordinates \\wideru\\\\ $^b$~Updated values of the proper motion (A.~Lommen 2001, private communications) \\ \\ \\ \\ \\ \\ $\\!$ $^e$~Dispersion measure \\\\ $^c$~Numbers in parentheses are uncertainties referring to the last significant digit quoted \\\\ $^f$~$D\\!M$-based distance (the new distance model by \\cite{cl02} places the pulsar to 317 pc) \\\\ \\end{table*}  Despite these differences, the distribution of integrated radio luminosities, as well as the luminosity dependence on $P$, $\\dot{P}$, $B$, and spindown energy losses $\\dot{E}$, are apparently similar for these much older and low-magnetized NSs, and for ordinary pulsars (\\cite{KL01}). About a dozen radio MSPs have been detected in X-rays. It is remarkable that their efficiency in converting spindown energy to X-ray luminosity is roughly the same as for ordinary pulsars, $L_{\\rm X}/\\dot{E}\\sim 10^{-3}$ (\\cite{BT97}; \\cite{Beck00}). This suggests  that the emission mechanisms responsible for the multi-wavelength radiation of MSPs and  ordinary pulsars  can be similar, and one could therefore expect to detect MSPs in other spectral ranges as well, as has been done for several ordinary pulsars. Detection of the first MSP in gamma-rays (\\cite{K00}) supports this idea.  To our knowledge, there are still no reports on optical detection of isolated MSPs. It is hardly possible to detect thermal emission from the entire surface of these old, $10^{8}-10^{10}$ yr, and cold NSs.  However, the spindown energy, expected to power the nonthermal emission of pulsars, can be much higher for MSPs than for old ordinary pulsars, and may even rival that of  young pulsars. Assuming the same efficiency of conversion of spindown energy to nonthermal optical luminosity as for ordinary pulsars, one can estimate that nearby MSPs may well be detectable in the optical with large telescopes. A problem is, however, that most of the nearby and energetic MSPs are components of close binary systems where the companion is predominantly either  a white dwarf or main sequence star (\\cite{Lor01}) which outshines the pulsar in the optical. Fortunately, there are at least nine {\\em solitary} MSPs in Galactic disk whose companions are believed to have been either evaporated or ablated (see, e.g., Lommen et al.\\ \\cite{L00}).  Here we report on deep BVR imaging of the field of one of these solitary millisecond pulsars, PSR J0030+0451. This pulsar was only recently discovered with the Arecibo telescope (Lommen et al.\\ \\cite{L00}), and soon thereafter detected in X-rays during the final observations with the {\\sl ROSAT/PSPC}  (\\cite{Beck00}). Recently it was re-observed in X-rays with the {\\sl XMM-Newton} (\\cite{BA02}). This relatively nearby NS (see Table~\\ref{t:param} for its parameters) is characterized by high X-ray flux, about $(4.3-6.8)\\times10^{-13}$ erg s$^{-1}$ cm$^{-2}$ in the $0.1-2.4$~keV band, and low interstellar absorption, $N_{\\rm H}\\la 3\\times10^{20} $ cm$^{-2}$, corresponding to a color excess $E(B-V)\\la 0.06$ mag. This makes it a promising candidate for optical detection. In Sect.~2 we present the observations and the data reduction. In Sect.~3 we discuss our results in the optical in conjunction with the available X-ray data. ",
        "conclusions": "\\label{s:disc}  \\begin{figure*}[t] \\includegraphics[width=130mm, angle=-90, clip]{h4067sp.eps} \\caption{ Radio and X-ray observations of \\psr\\ with the Arecibo telescope (Lommen et al.\\ \\cite{L00}), {\\sl ROSAT} (\\cite{Beck00}), and {\\sl XMM-Newton} (\\cite{BA02}), as well as VLT upper limits in the BVR${\\rm_s}$ optical bands. Solid, dashed, and dot-dashed lines show the best spectral fits of the  {\\sl XMM-Newton} data with three alternative two component  spectral models: BB+PL, broken PL, and BB+BB, respectively (BB stands for blackbody). Parameters of the fits and a mean unabsorbed integrated flux $F_{\\rm X}$, approximately the same for all the fits (\\cite{BA02}), are indicated in the plot. All fits are acceptable at the same significance level in the $\\sim$ $0.3-7$~keV range, indicated by a horizontal bar with arrows at the upper-left, and  their $\\sim$90\\% uncertainties are shown as shaded regions. The unabsorbed model spectra are extrapolated toward the optical range. Previous {\\sl ROSAT} data contain no spectral information and provide only an integrated flux indicated in the figure by a bold cross (\\cite{Beck00}), roughly consistent with the more recent {\\sl XMM-Newton} results. The VLT upper limits, marked by end bars, show  that any nonthermal PL  component obtained from the above spectral fits of the X-ray data must be strongly suppressed in the optical range. This implies  either a cutoff or a strong break in the $0.003-0.1$~keV range. The purely thermal BB+BB model is consistent with our optical upper limits. Further details of this plot are discussed in Sect.~\\ref{s:disc}. } \\label{f:flux} \\end{figure*}  Based on the apparent similarity  of the radiation properties of MSPs and ordinary pulsars, and assuming the same mechanisms to generate optical emission in both types of pulsar, we expected to detect the optical counterpart of \\psr\\ at the $m\\sim26$ visual magnitude level, assuming a simple scaling of $\\dot{E}/d^2$ from known optical fluxes of ordinary pulsars. Our observations were deeper, but still we did not detect any reliable counterpart candidate.  To try to understand what our optical non-detection implies, we have plotted the available information about the multiwavelength spectrum of \\psr\\ including radio, optical, and X-ray data in Fig.~\\ref{f:flux}. For the X-ray region we have included preliminary results of recent {\\sl XMM-Newton} observations (\\cite{BA02}) where the pulsar was clearly detected in the $0.3-7$~keV range. The detection range is shown by a horizontal bar with arrows in Fig.~\\ref{f:flux}. The overall spectrum compares well with multiwavelength spectra of ordinary pulsars (e.g., \\cite{Kop01}), i.e., the pulsar flux is higher in the radio and fades toward the X-ray range. This can be explained by different emission mechanisms in the radio (coherent) and at shorter wavelengths (non-coherent). However, a more detailed inspection of the optical/X-ray range reveals a feature which has not been seen for ordinary pulsars. For the latter, the optical flux is usually close to an extrapolation of the nonthermal high energy tail of the  X-ray emission, usually described by a power-law (PL). From Fig.~\\ref{f:flux} it is clear that this is not the case for the \\psr.  Fig.~\\ref{f:flux} shows that in the {\\sl XMM-Newton}  range the data can be fitted equally well by three different two-component spectral models: a blackbody + power-law model (BB+PL), a broken (or curved) PL model, or a model based on two different blackbodies (BB+BB) with $N_{\\rm H} \\le 2.5\\times10^{20}$ cm$^{-2}$. From the X-ray data alone it is difficult to discriminate between the three models, although the sharp X-ray pulse profile perhaps favors the domination of a nonthermal PL component (\\cite{BA02}). The VLT upper limits allow additional constraints using an extrapolation of the X-ray model spectra toward the optical range.    From Fig.~\\ref{f:flux} it is seen that the low-energy extension of the BB+PL fit overshoots the optical flux upper limits of \\psr\\ by $\\sim$5 orders of magnitude. The broken PL extension, which implies a flatter spectrum in the soft X-ray energy band, is also 3-4 orders of magnitude higher. We have not corrected the data for interstellar extinction, which is low and does not play any significant role at such large differences. The strong suppression of the PL components in the optical range suggests  that these two models either should be ruled out, or that their PL components must  have a strong break or even a cutoff at a photon energy somewhere in the $0.003-0.1$~keV range. Such a situation has never been observed for any ordinary pulsar. For example, the spectral index $\\alpha_{\\nu}$ (defined as $F_{\\nu} \\propto \\nu^{-\\alpha_{\\nu}}$) of the young Crab pulsar changes from $\\sim0.5$ for soft X-rays to zero in the FUV/optical/near-IR range (e.g., \\cite{S00}; \\cite{S02}). For the relatively young Vela-pulsar (\\cite{MC01}), the middle-aged PSR B0656+14 (\\cite{Kop01}) and PSR B1055$-$52 (\\cite{P02}), and even for the old ordinary pulsars PSR B0950+08 (\\cite{Zhar02}) and PSR B1929+10 (\\cite{M02}), the extension of the PL X-ray component matches the optical flux, suggesting the same mechanism of nonthermal emission in the optical and X-ray ranges. It is not clear what could be the reason for such a strong change in the nonthermal spectral slope of \\psr\\ from negative in X-rays to positive in the optical range. As seen from Fig.~\\ref{f:flux}, the X-ray PL fits and our upper limits exclude any flat extension toward the optical range, even if we place a break point energy at the lower boundary of the {\\sl XMM-Newton} range $\\sim 0.1$~keV. Note that the low-frequency extensions of the PL components overshoot even the radio fluxes. This is also not typical for ordinary pulsars. One can  assume that nonthermal emission is  due to  synchrotron  radiation of relativistic particles in the magnetosphere of the pulsar. In this case we obtain for the simplest monochromatic particle distribution over the energy a  spectral flux of $F_{\\nu} \\propto \\nu^{1/3}$ in the low frequency range below a maximum frequency $\\nu_{\\rm m} \\propto B\\gamma^2$. Here $B$ is the magnetic field, and $\\gamma$ is the gamma-factor of the emitting particles. From the spectral shape suggested by the X-ray data and optical upper limits it is natural to put $\\nu_{\\rm m}$ near the maximum of the X-ray spectral flux at the low boundary of the {\\sl XMM-Newton} range. At typical gamma-factors of primary and secondary relativistic particles in pulsar magnetospheres, i.e. $\\sim 10^6$ and $\\sim 10$, respectively, and  for the period of \\psr\\ $\\sim4.9$~ms, the same peak frequency value is predicted by a model of synchrotron emission from the pulsar light cylinder suggested  by \\cite{Mal01}. For the expected synchrotron flux below the adopted $\\nu_{\\rm m}$ value (see Fig.~\\ref{f:flux}, dotted lines) we would likely hint the optical counterpart, but we did not.    The purely thermal BB+BB spectral model is consistent with our upper limits without any additional assumptions. Its Rayleigh-Jeans tail is about 6 stellar magnitudes fainter in the optical than our upper limits, and would hardly be detectable with present telescopes. Thermal photons can be emitted by hot polar caps of the pulsar. The two-blackbody fit  indicates a non-uniformity in the temperature distribution over the caps. This could be described by heat propagation over the surface of the neutron star out of a hot cap core, including neutron star atmosphere effects, as has been done in the case of the MSP J0437$-$4715 (\\cite{Z01}). An additional faint PL component is required to fit an excess over the thermal emission at high energy  X-rays from PSR J0437$-$4715.  If the same would be true for \\psr, it could be brighter in the optical  than estimated from the simple BB+BB model due to a contribution from a similar nonthermal component of magnetospheric origin. Deeper X-ray observations are probably needed to detect this component  in the high energy tail of the  \\psr\\ spectrum. The similarity of X-ray and radio pulse profiles of this pulsar suggests that radio and X-ray peaks are in phase (\\cite{BA02}), although direct timing to confirm this has not yet been done. In the frame work of the thermal model this means that radio emission is generated close to the polar cap surface and the similarity of the pulse shapes is likely caused by the same geometry of the emitting regions.   Based on our upper limits we can constrain also the efficiency of converting spindown power of  \\psr\\ to optical emission. The luminosity in the B band is $L_{\\rm B}=4\\pi d^2F_{\\rm B} \\Delta\\nu_{\\rm B} \\la 6.3 \\times 10^{26}~d_{230}^2$ erg s$^{-1}$, and hence the optical efficiency $\\eta_{\\rm B} =L_{\\rm B} /\\dot{E} \\la 1.9\\times 10^{-7} d_{230}^2$. Here $d_{230}=d/230$~pc is the normalized distance to \\psr. This upper limit is about half the efficiency of the middle-aged PSR~B0656+14 (at 500 pc) and exceeds the efficiencies of the Geminga and Vela pulsars by about 1 and 2 orders of magnitude, respectively (see, e.g., \\cite{Zhar02}). It is interesting to note that the expected efficiency in the BB+BB model in Fig.~\\ref{f:flux} is 2-3 orders of magnitude lower than the upper limit derived above from our optical data, and that the efficiency in the BB+BB model is comparable to that of the Vela pulsar.  The comparison of the efficiencies makes sense here, since the thermal emission from hot polar caps and the nonthermal magnetospheric emission of the Vela-pulsar, both are powered by relativistic particles produced in magnetospheres of rapidly rotating NSs.  Thus, we see that the optical efficiency of the \\psr, as derived from our and X-ray observations, is not unusually low, but is compatible with the efficiency range of ordinary pulsars detected in the optical band.  It has been shown for ordinary pulsars that spectral index appears to became steeper with pulsar age in the optical range (\\cite{Kop01}; \\cite{MC01}), while it flattens in gamma rays (e.g., \\cite{sg02}). It has been noticed also across a restricted set of young and middle-aged pulsars detected in the optical and gamma regions that the gamma-ray efficiency increases with age, while the reverse is true for the optical efficiency (\\cite{gold95}). This would suggest that there is a reprocessing of the gamma-photons into the optical in pulsar magnetospheres, and that it is more efficient for younger pulsars than for older ones (e.g., \\cite{sg02}). Thus, it would be not surprising that very old \\psr\\ is fainter in the optical than we expected. However, recent optical studies of old ordinary pulsars (\\cite{Zhar02}; \\cite{M02}; \\cite{M02b}) have revealed nonmonotonous behavior of the optical efficiency {\\it vs} age with a minimum at  $\\tau \\sim 10^4-10^5$~yr and further increase towards higher ages $\\tau \\ga 10^7$~yr. Old pulsars can be actually much more efficient than the middle-aged ones and produce the optical photons with almost the same efficiency as young and energetic Crab-like pulsars. In this context low optical brightness of the MSP J0030$+$0451 remains puzzling.   A clue to what dominates the X-ray and optical spectrum would hopefully be found from observations of \\psr\\ and other MSPs in the FUV and, in particular, in the EUV range. Even deep upper limits in these ranges would help to understand how strongly the multiwavelength emission and radiation properties of MSPs differ from those of ordinary pulsars."
    },
    "0212/astro-ph0212392_arXiv.txt": {
        "abstract": "The Sunyaev-Zel'dovich (SZ) effect of galaxy clusters is characterized by three parameters: Compton parameter, electron temperature and cluster peculiar velocity. In the present study we consider the problem of extracting these parameters using multi-frequency SZ observations only. We show that there exists a parameter degeneracy which can be broken with an appropriate choice of frequencies. As a result we discuss the optimal choice of observing frequencies from a theoretical point of view. Finally, we analyze the systematic errors (of the order $\\mu$K) on the SZ measurement introduced by finite bandwidths, and suggest a possible method of reducing these errors. ",
        "introduction": "The interaction of the Cosmic Microwave Background Radiation (CMBR) photons with the free electrons in the ionized gas of clusters of galaxies produces a small change in the intensity known as the Sunyaev-Zel'dovich (SZ) effect. This distortion arises from the transfer of photons from the low-energy to the high-energy or Wien side of the planckian spectrum. The underlying physics of the SZ effect is well understood and a quantitative description was given by Sunyaev \\& Zel'dovich \\cite{sz}, who already realized its cosmological significance. Nowadays there exists reliable observations of the SZ effect (see e.g.~\\cite{LaRoque01,coma,joy2001}) and dedicated experiments are being prepared, which will provide information about the kinematics and evolution of clusters, which in turn can be used to extract cosmological parameters (see e.g.~\\cite{rep95,bir99,car02} for recent reviews).  The intensity change of the CMBR caused by the SZ effect is proportional to the Comptonization parameter, \\begin{equation} y_c = \\int dl~ \\frac{T_e}{m_e} ~ n_e \\sigma_{\\rm Th} ~, \\end{equation} where $T_e$ is the temperature of the electron gas in the cluster, $m_e$ the electron mass, $n_e$ the electron number density, $\\sigma_{\\rm Th}$ the Thomson scattering cross section, and the integral is calculated along the line of sight through the cluster. We use units for which $k_B=\\hbar=c=1$.  For an isothermal intra-cluster gas one has $y_c=\\tau \\, T_e/m_e$, where $\\tau=\\int dl~ n_e \\sigma_{\\rm Th}$ is the optical depth. The distortion of the CMBR is then given by an intensity change \\begin{equation} \\Delta I_{\\rm T} = I_0~ y_c~ \\left( f_0(x)~ +\\delta f(x,T_e)\\right)~, \\label{thermal} \\end{equation} where \\begin{equation} f_0(x) = \\frac{x^4e^x}{(e^x-1)^2}\\left[ \\frac{x(e^x+1)}{e^x-1}-4\\right]~. \\label{f0} \\end{equation}  \\begin{figure} \\begin{center} \\epsfxsize=9cm \\epsffile{fig1.eps} \\end{center} \\caption{The normalized intensity change caused by the thermal and kinetic SZ effects for the choice $T_e=15$ keV and $v_p=500$ km/sec. The solid line is the non-relativistic SZ effect. The dashed lines are the relativistic corrections, both alone and together with the thermal effect.  The dot-dashed lines are the kinetic SZ effect, both alone and together with the thermal effect.} \\label{fig:example} \\end{figure}    Here $x=\\nu/T_0$ is the dimensionless frequency with $T_0 =2.725$ K, and $I_0=T_0^3/(2\\pi^2)$. The intensity change is independent of the temperature for non-relativistic electrons, a limit which is valid for small frequencies ($\\nu\\simlt 100$ GHz), but for high frequencies it must be corrected \\cite{rep95b} with $\\delta f(x,T_e)$ either by using an expansion in $\\theta_e=T_e/m_e$ \\cite{cha98,ito98,noz00} or calculated exactly \\cite{dol01}. In some rich clusters the optical depth is large enough ($\\tau \\sim 0.02-0.03$) to require the inclusion of multiple scattering effects \\cite{dol01,ito01,colaf}.   When the cluster has a relative motion with respect to the CMBR rest frame, there exists an additional distortion known as the kinetic SZ effect \\cite{sz80}, which is identical to a change in the temperature of the CMBR.  The corresponding intensity change, for a peculiar velocity $v_p$, is $\\Delta I_{\\rm K} = I_0 y_c f_{\\rm kin}$, where \\begin{equation} f_{\\rm kin} = -   v_p \\frac{m_e}{T_e} \\frac{x^4 e^x}{(e^x-1)^2}~. \\label{kinematic} \\end{equation} This expression is slightly modified by the corrections caused by the electron temperature \\cite{saz98,noz98}.  The total intensity change is just the sum of the thermal and kinetic SZ effects, \\begin{equation} \\Delta I_{\\rm total} = \\Delta I_{\\rm T}+\\Delta I_{\\rm K}~, \\label{total} \\end{equation} and it is characterized by three parameters: $y_c, T_e$ and $v_p$. The different contributions to the SZ distortion are presented in \\fref{fig:example}, where the total intensity change in \\eref{total} is compared to the thermal effect with vanishing temperature (the first term in \\eref{thermal}). The kinetic SZ effect and the contribution of the relativistic corrections (the second term in \\eref{thermal}) are presented for the parameter choice $T_e=15$ keV and $v_p = 500 ~{\\rm km}/{\\rm s}$.  Let us emphasise that all 3 SZ parameters can in principle be extracted from multi-frequency observations of the SZ effect. This is because the frequency dependence of the SZ effect is different for the different parameters. This is clear from \\fref{fig:example}, where one can see that the main effect (the non-relativistic SZ effect characterized by the comptonization parameter, $y_c$) has one crossing with the zero-line. The kinetic effect (governed by $v_p$) has no crossings of the zero-line, and the relativistic effect (governed by $T_e$) crosses the zero-line twice. Thus with several observing frequencies placed optimally and with good sensitivity, one can extract all 3 SZ parameters as we describe in detail in \\sref{optimal}.  In this paper we discuss the possibilities for future SZ observations to extract these 3 cluster parameters, taking into account the parameter degeneracies described in \\sref{sec:deg}. To this end we try to extract the 3 SZ parameters from a set of simulated galaxy clusters with realistic temperatures and peculiar velocities, using the sensitivities of upcoming experiments like ACT and Planck. In \\sref{optimal}, we discuss the optimal observing frequencies from a theoretical point of view. Finally, we analyze in \\sref{bandwidth} the systematic error introduced by a finite frequency bandwidth, which may be as large as few $\\mu$K and hence important for the next generation of SZ observations. ",
        "conclusions": "We have pointed out that there are parameter degeneracies for the Sunyaev-Zel'dovich effect. This in particular implies that the error-bars on the peculiar velocities for clusters with positive peculiar velocities is significantly smaller than for cluster with negative peculiar velocities. This parameter degeneracy can be broken either by having prior temperature knowledge, or by observing at several frequencies. For an experiment with $2~\\mu$K sensitivity and 3 observing frequencies (ACT-like), we find that a Compton parameter of the order of $10^{-4}$ is often deduced with a good accuracy 2\\% (10\\%) with (without) prior information on the electron temperature. The prior information on the electron temperature is even more important for the extraction of the peculiar velocity for which the accuracy reaches 10--20\\%. Additional observing frequencies also break the parameter degeneracies but the accuracies are not better than with the prior temperature knowledge. In the particular case of Planck, the Compton parameter of $~10^{-4}$ is extracted very accurately (2\\%) even without information on the temperature. Prior temperature knowledge is, however, necessary to go from a peculiar velocity determination of a factor of a few to 20--50\\%.  We have discussed the needed sensitivity in order to extract the cluster temperature using SZ observations only. We have seen that experiments with only 3 observing frequencies need very good sensitivity ($\\sim 0.5 ~\\mu$K) in order to determine the cluster temperature with 1 keV error-bars. Already with 4 frequencies one can reach 1 keV (4 keV) temperature error-bars with only $1~\\mu$K ($4~\\mu$K) sensitivity.   We have identified the frequencies which are optimal for deducing the 3 cluster parameters from a theoretical point of view.  This optimal choice of 4 observing frequencies and sensitivity of 1 $\\mu$K allows one to deduce the parameters of a 5 keV, $y_c=3\\times 10^{-5}$ cluster with $1\\sigma$ error bars of about 2 keV and $10^{-6}$. The velocity is extracted with a 50\\% accuracy.  The optimal frequencies for an actual experiment will depend both on the sensitivity of the experiment and also on which cluster parameters are of primary interest.  We have studied the systematic error which appears due to the finite band-width.  This error is related to the shape of the SZ effect. The exact shape of the SZ effect depends on the cluster temperature and peculiar velocity, which cannot be known a priori. This systematic error can be of the order of a few $\\mu$K, and must therefore be considered seriously by future experiments which aim at such impressive sensitivity.  \\ack It is a pleasure to thank Lloyd Knox, Francesco Melchiorri, Fran\\c cois Pajot and Rocco Piffaretti for useful discussions. The work of DS was partly supported by the Deut\\-sche For\\-schungs\\-ge\\-mein\\-schaft under grant No.\\ SFB 375. SP was supported by the Spanish grant BFM2002-00345 and a Marie Curie fellowship under contract HPMFCT-2002-01831.   \\label{lastpage}"
    },
    "0212/astro-ph0212321_arXiv.txt": {
        "abstract": "The study of both neutral and ionized gas in young radio sources is providing key information on the effect the radio plasma has on the ISM of these objects.  We present results obtained for the compact radio sources PKS~1549--79, 4C~12.50 and PKS~1814--63 and for the intermediate-size radio galaxy 3C~459.  At least in the first two, low ionisation optical emission lines and HI absorption appear to be associated with the extended, but relatively quiescent, dusty cocoon surrounding the nucleus.  The [OIII] lines are, on the other hand, mostly associated with the region of interaction between the radio plasma and the ISM, indicating a fast outflow from the canter.  A case of fast outflow (up to $\\sim 1000$ \\kms) is also observed in HI in the radio source 4C~12.50. As the radio source evolves, any obscuring material along the radio axis is swept aside until, eventually, cavities (of the same kind as observed e.g.\\ in Cygnus ~A) are hollowed out on either side of the nucleus. We may witness this phase in the evolution of a radio source in the radio galaxy 3C~459. ",
        "introduction": "Although in the recent years important progress has been made in the understanding of the physics of AGN, a number of main questions still remain open. One of them is the early evolution of powerful radio galaxies.  To know more about this phase is important not only from the point of view of the detailed phenomenology of these objects, but also for understanding the evolution of massive galaxies.  In general, galaxies in their early stage of radio activity are likely to have their nuclear regions enshrouded in a cocoon of material left over from the accretion event(s) that trigger the activity. As the activity evolves, and the radio jets expand, the ISM and dust along the radio axis will be swept from the nuclear regions by jet-cloud interactions and/or quasar-induced winds. Thus, outflows of gas can be significant in this phase.  These processes have been suggested to profoundly affect even the star formation history in luminous galaxies (Silk \\& Rees 1998). Thus, the study of the early-phase of radio activity and its effect on the galaxy medium has broad implications.  As described in many contributions in these Proceedings, both neutral and ionized gas are clearly present around young radio sources.  The study of this gas has a key role in the investigation of the evolutionary scenario described above. Spectral line observations can, in fact, provide information on the gas kinematics where signatures of interaction with the radio plasma can be found.  Here we present the results (some still preliminary) obtained in the study of the HI in three compact radio sources (PKS~1549--79, 4C12.50 and PKS~1814--63) and one intermediate-size radio galaxy (3C~459).  While the HI detected in compact radio sources is very often interpreted as due to a circum-nuclear disk/torus, it is now clear that more complicated situations are present (see e.g. Morganti 2002 and ref. therein) and that neutral gas can also trace the presence of (sometimes extreme) interactions.  The objects described here are also part of a detailed study of their ionized gas. This information is crucial for the correct interpretation of the kinematics of the neutral hydrogen, as it will be illustrated below.   \\begin{figure} \\centerline{\\psfig{file=Morganti.fig1.ps,height=7cm} \\psdraft \\psfig{file=Morganti.fig1.ps,height=7cm} \\psfull} \\caption{{\\sl Left:} profile of the HI absorption in PKS~1549--79 (from Morganti et al. 2001). The velocities derived from the optical emission lines are indicated. {\\sl Right:} continuum radio emission on the VLBI scale from King 1996.} \\end{figure} ",
        "conclusions": "We have shown as (at least some) compact radio galaxies can have different gas components with very different kinematics.  The neutral hydrogen does not appear to be associate only to circum-nuclear disk/tori (as often claimed), but instead there is evidence of the presence of more diffuse and quiescent halos in which these young sources are embedded.  The radio jets are expected to clear their way through the rich ISM.  As indication of this, high speed outflow are detected not only in the ionized gas but, quite surprisingly, also in neutral hydrogen in the case of 4C~12.50.  The kinematics of the neutral hydrogen alone can be sometimes difficult to be unambiguously interpreted. However,  it can provide crucial information on the condition, morphology and kinematics of the gas in the central regions of radio sources once it is combined with the results from the ionized gas.  This has been clearly shown for the case of the third radio galaxy presented, PKS 1814-63.  Despite the complex and intriguing HI absorption detected, the interpretation is still uncertain due to the lack of high quality optical spectra.  Finally, we have presented a case of intermediate-size radio source (3C~459) where we may witness the next step of the evolution of the source.  In how many young radio sources do we expect to see the effect of the radio plasma making its way through the ISM and producing gas outflows?  We do not have yet enough information available, especially in the optical (e.g.\\ accurate measurement of the redshift), to estimate this in a reliable way.  The group of radio galaxies that are far-IR bright and that show a young stellar population in their optical spectra seems, however, to be the most likely to show these extreme effects and outflows."
    },
    "0212/astro-ph0212117_arXiv.txt": {
        "abstract": "The limiting factor of the accuracy of WFPC2 photometry is the CTE loss, which has increased to the level of 50\\% or more for faint stars at the top of the chips.  I describe recent work on characterizing this effect, and provide improved equations for CTE correction.  I also examine issues affecting background measurement, which if not done correctly can introduce artificial nonlinearities into photometry. ",
        "introduction": "Several obstacles inhibit the obtaining of accurate photometry from WFPC2 images.  The most severe of these is charge loss during readout, commonly known as CTE loss.  While no CCD is likely to be perfect, the effect is pronounced on WFPC2, where initial measurements showed that a star at the top of the chip ($y = 800$) would lose approximately 10\\% of its charge while being read out.  This loss was reduced by cooling the camera from $-76^{\\circ}$C to $-88^{\\circ}$C, and Holtzman et al. (1995) found that the CTE loss could be corrected to acceptable levels with a correction of $0.04\\ y/800$ magnitudes to photometry.  Unfortunately, this has increased over time, and faint stars at the top of the chip now can lose well over half their light to CTE loss.  Improved characterizations of the charge loss were produced by Stetson (1998), Whitmore, Heyer, \\& Casertano (1999), and Saha, Labhardt, \\& Prosser (2000) by analysis of larger data sets using DAOPHOT, IRAF apphot, and DoPHOT, respectively.  While there were several issues of agreement between these studies, there were also significant discrepancies.  For example, Stetson (1998) found no significant time dependence, while Whitmore et al. (1999) did.  Likewise, Saha et al. (2000) found no X-CTE loss and no nonlinearity, while both were observed in the other studies.  The fact that a different photometry package was used in each study raised the possibility that the observed CTE loss was partly a function of the package used.  Another possible source of photometric error was reported from measurements of the same fields in long and short exposures, in which objects appeared fainter in the short exposure than in the long exposure.  This was independent of the position on the chip, and the phenomenon has become known as the ``long vs. short anomaly''.  Casertano \\& Mutchler (1998) characterized this effect using short and long observations of NGC 2419, and found it to be a function of counts rather than exposure time.  The effect they found was a very large one; rather than the 5\\% effect reported previously, their correction is 0.18 magnitudes at 280 electrons (50 ADU at gain of 7).  However, Stetson (1998) solved for a position-independent effect as part of his CTE study and found none.  Given that his data set included the NGC 2419 observations, this only added to the questions about how well these effects were really understood.  Dolphin (2000a; hereafter D00) presented a study of the WFPC2 CTE based on reductions of 843 WFPC2 images of $\\omega$ Cen and NGC 2419 using his HSTphot photometry package (Dolphin 2000b).  In addition to providing yet another CTE solution based on yet another photometry package, this work shed light on the discrepancies noted above.  The lack of time dependence seen by Stetson was observed to be primarily the result of an insufficient time baseline and the fact that all of his high-background data were obtained at the end of the time baseline.  The lack of a nonlinearity in the study of Saha et al. was found to result from their background measurements, which contained significant amounts of starlight.  This meant that their CTE correction (explicitly a function of background only) was implicitly a function of brightness as well.  Finally, it was demonstrated that the long-short anomaly is primarily a result of poor background measurements, as my reductions of the NGC 2419 data showed no significant discrepancies after the CTE loss was corrected.  In this paper, I describe ongoing efforts to improve upon D00 and present the latest results. ",
        "conclusions": "I have developed and described a CTE correction procedure to supersede that of D00.  There are several significant improvements over the previous corrections.  First, a single uniform set of ground-based photometry is used to provide the comparisons.  Second, HSTphot has been improved several times over the intervening time.  Third, more stringent cuts were applied to the data to eliminate bad points.  Finally, an improved functional form of the CTE correction accounts for the changing CTE loss as a star becomes fainter during readout.  I have also examined a commonly-overlooked aspect of stellar photometry, background measurement.  In short images, the background value is sufficiently low that the digitization effects dominate the histogram.  If not handled correctly, this can result in significant artificial nonlinearities added to the data.  After exploring the options offered by IRAF, the recommendation is that one use the ``mean'' algorithm with sigma clipping.  Finally, I conclude by describing ongoing efforts to improve the CTE corrections for faint sources and to obtain an image-based CTE correction.  My WFPC2 calibration web site (http://www.noao.edu/staff/dolphin/wfpc2\\_calib/) will be kept updated with improvements to the CTE corrections and photometric zero points as they become available."
    },
    "0212/astro-ph0212267_arXiv.txt": {
        "abstract": " ",
        "introduction": "Gamma-ray bursts observed with space-based detectors span a wide range of durations limited by the detector integration time to time scales $\\rm t \\ge 1ms$. Satellite-based detectors are not efficiently triggered by pulses with time scales shorter than $\\rm \\sim 1ms$ because their effective collection area is too small to allow for faster time sampling of signals with typical fluences. Nevertheless, among the detected bursts, millisecond and sub-millisecond variability is common (Walker \\& Schaefer 2000). \\begin{figure}[t] \\begin{center} \\includegraphics[height=17pc]{pbhtechimg.eps} \\end{center} \\caption{In an imaging telescope, $\\gamma$-ray bursts should appear as a Cherenkov glows.} \\label{principle} \\end{figure}  A front of low energy $\\gamma$-rays entering the Earth's atmosphere results in the development of small electromagnetic cascades at $\\rm \\sim 20km$ altitude. Each of them produces Cherenkov light spread over an area of $\\rm \\sim 500m$ in radius at ground level. The Cherenkov light from all the showers should give a detectable glow extending over $\\rm \\sim2^o$ centered on the direction of the burst (figure~\\ref{principle}). The diffuse night sky background (our sensitivity limitation) of collected over such a solid angle during $\\rm \\sim10\\mu s$ is comparable to the Cherenkov light yield from a few $\\rm 100MeV$ $\\gamma$-ray burst with a fluence of $\\rm 10$ $\\gamma$-ray $\\rm \\cdot m^{-2}$. This also corresponds to the fluence sensitivity of a space-based detector with collection areas of $\\rm \\sim 1 m^2$. Going toward shorter time scales, the noise contamination is reduced and atmospheric Cherenkov detectors gain in sensitivity following the square root of the burst duration. The fluence sensitivity for a 100~ns burst of 1~GeV $\\gamma$-rays is $\\rm \\sim 0.1 \\gamma ray \\cdot m^{-2}$ (Krennrich et al., 2000).  The timing and imaging analysis allows to remove the dominant $\\rm 100ns-10\\mu s$ background due to atmospheric scintillation produced by ultra-high-energy cosmic-ray showers. The features of the multi-$\\gamma$-ray front initiated shower images are unique. Both the light distribution and the time structure of the Cherenkov light image are centered on the direction of the burst and should be symmetric. We are not aware of any phenomenon capable of producing fake signals and we expect the technique to allow background free detection of bursts.  SGARFACE is the first ground-based experiment using the imaging and time sampling to search for ultra-short bursts of $\\gamma$-rays. The Whipple observatory 10~m $\\gamma$-ray telescope will be used to collect the Cherenkov light and form the image. After discussing the design and status of SGARFACE, we will review some of its motivations. ",
        "conclusions": "With the installation of SGARFACE at the Whipple 10m telescope, its sensitivity will be expanded to GeV photon bursts of $\\rm 0.1\\mu s$ to $\\rm 100\\mu s$.  SGARFACE is expected to have a fluence sensitivity of 500 times that of EGRET for microsecond bursts of GeV photons. Data taking should start before the end of 2002 and will be the object of an interesting search for sub-millisecond $\\gamma$-ray signal with possible implications for primordial black holes and pulsar giant pulses."
    },
    "0212/astro-ph0212098_arXiv.txt": {
        "abstract": "{We have determined Li abundances for 11 G--type stars in the 6--8~Gyr old open cluster \\object{NGC~188}. These data significantly enlarge the number of cluster stars with Li measurements, allowing us to extend the investigation of Li depletion in open clusters to ages well beyond the age of the Sun. We have also inferred the cluster metallicity which turns out to be solar. We find that solar--type stars in NGC~188 are only slightly more Li depleted than the much younger \\object{Hyades} and no more Li depleted than stars of similar temperature in the 2--4~Gyr younger cluster \\object{M~67}. At variance with M~67, NGC~188 members show virtually no scatter in their Li abundances. Surprisingly,  no solar--type star in NGC~188 appears as Li depleted as the Sun or as the most Li depleted stars in M~67. We discuss the implications of these results for mechanisms of internal mixing and Li depletion in main sequence stars. ",
        "introduction": "Observations of light elements in stars are powerful tools to test Big Bang nucleosynthesis, Galactic chemical evolution and stellar interior mixing processes. Lithium has a very low burning temperature ($\\mathrm{\\sim 2.5\\times 10^{6}\\,K}$) and it survives only in the outermost layers of a star; therefore, it is a good tracer of mixing mechanisms and stellar interior structure. Observations of Li in clusters of different ages and chemical compositions allow us to put constraints on these processes and to investigate their dependence on age, metallicity and possibly other stellar parameters. Observations of cluster stars carried out during the past 15 years have shown that Li depletion is not a simple function of age and metallicity as predicted by standard models (see e.g. the reviews of Jeffries \\cite{jef00} and Pasquini \\cite{pas00} and references therein). For example, standard models, which include only convection as a mixing process, predict that little (if any) Li depletion should occur during the main sequence (MS) evolution of solar--type stars, since the convective zone is too shallow to reach the Li burning layer; the comparison of the Li abundance distributions of clusters of different ages (spanning the range from the zero age main sequence -ZAMS- to the solar age) indicates instead that solar--type stars do destroy Li on the MS. \\begin{figure} \\psfig{figure=MS2971f1.eps, width=11.0cm, angle=0} \\vspace{-2cm} \\caption{Sample spectra in the Li region.} \\end{figure} Furthermore, the solar age, solar-metallicity cluster M~67 is characterized by a large dispersion in Li abundances for stars of similar temperatures (Spite et al. \\cite{spi87}; Garc\\'\\i a L\\'opez et al. \\cite{gar88}; Pasquini et al. \\cite{pas97}; Jones et al. \\cite {jon99}); about 60 \\% of the stars appear Li-rich, with a Li abundance a factor of 2--3 below the 600~Myr old Hyades, while the remaining fraction are Li--poor with a factor of 5--10 lower abundances. This result shows that additional parameters besides mass, age, and metallicity affect Li depletion, and that non-standard mixing processes such as diffusion, mass loss, or rotationally driven mixing, are at work.  The comparison between M~67 and the intermediate age clusters \\object{IC~4651}, \\object{NGC~3680}, and \\object{NGC~752} (age $\\sim 2$~Gyr) shows instead that stars in the upper envelope of M~67 have similar abundances as the intermediate cluster stars, suggesting that either Li depletion becomes very slow at old ages, at least for a fraction of stars, or that different conditions (e.g., heavy element abundances, initial angular momentum distributions and/or rotational evolutions) in different clusters lead to different Li evolutions (Randich et al. \\cite{r00}). We also mention that neither IC~4651 nor NGC~3680 show any scatter in Li, while a scatter might be present among NGC~752 members cooler than 5800~K (Hobbs \\& Pilachowski \\cite{hp86}).  In order to put additional empirical constraints on MS Li depletion, in particular on its timescale and on the development of the dispersion, we carried out a Li survey of the old open cluster NGC~188. This cluster has an estimated age of 6--8~Gyr (see discussion in Sarajedini et al. \\cite{sara99}, hereafter S99) and is one of the oldest known open clusters. According to S99 it is $3.0\\pm 0.7$~Gyr older than M~67. The only published data on Li abundances in dwarf stars in this cluster date back to nearly 15 yrs ago and were limited to only a few stars (Hobbs \\& Pilachowski \\cite{hp88}). A more detailed Li study of this cluster is thus warranted. Our paper is organized as follows: in Sect.~2 we describe our sample and the observations; the abundance analysis and the results are summarized in Sects.~3 and 4, while in Sects. 5 and 6 we present a discussion of the results and our conclusions. ",
        "conclusions": "We have determined Li abundances for 11 G--type stars (with dereddened B--V colors 0.57 $\\leq$ \\bmo $\\leq 0.65$) in the old open cluster NGC~188. Our data, together with the Li abundances of five stars (three in common with us) previously measured by Hobbs \\& Pilachowski (\\cite{hp88}), allow us to extend the investigation of Li depletion in Galactic open clusters to ages 2--4~Gyr older than the Sun. Two are the main results of this study:\\\\ {\\bf a)} solar--type stars in NGC~188 have all the same Li abundances (\\nli$=2.34 \\pm 0.13$) with virtually no scatter, contrary to what has been observed in M~67; {\\bf b)} solar--type stars in NGC~188 are no more Li depleted than stars in the upper envelope of the Li abundance vs. temperature distribution of M~67. We have also estimated the cluster metallicity finding [Fe/H]$\\simeq 0$.  The above results represent further challenges to our interpretation of Li depletion in MS solar--type stars. Based on the results for M~67 and field stars, one would have expected to find a significant fraction of Li-poor stars in NGC~188. Finding that 13 out of 13 members of NGC~188 have a Li abundance only slightly below the Hyades confirms that very old stars are not necessarily Li depleted; at the same time, the absence of NGC~188 members as Li depleted as the Sun or as the most depleted stars in M~67, appears intriguing and raises questions about the possibility of using a single open cluster as truly representative of all clusters of the same age and metallicity.  The comparison of NGC~188 with the upper envelope of M~67, with the 2~Gyr clusters, and with the Hyades shows that the average Li abundance is virtually the same for cluster with ages $\\geq 2$~Gyr. Whereas this result can be due to differences in heavy element abundances leading to different Li depletion timescales for different clusters, this possibility appears to be not very likely. We rather believe that the plateau in Li abundances observed for ages older than 2~Gyr supports the idea that, for part of the stars, Li depletion may become ineffective beyond this age. A suitable mixing process with these characteristics must be investigated. The spread among M~67 and in the field shows that part of the stars instead undergo a very dramatic additional Li destruction after 2~Gyr. However M~67 remains so far the only old cluster for which a dispersion among solar--type stars has been confirmed. The reasons for this remain unclear, but we suggest that M~67 might represent an inhomogeneous sample in some respect.  We finally mention that the average Li abundance of NGC~188 is only slightly higher than the Li plateau for Pop.~{\\sc ii} stars with [Fe/H] $\\leq -1.4$ in the halo and in globular clusters. Although this may be only a coincidence, it indicates the need of determining Li abundances in other old clusters (both open and globular) in order to link together the Li abundance of Pop.~{\\sc ii} stars (usually assumed to be equal to the primordial value) with that of Pop.~{\\sc i} stars in old open clusters."
    },
    "0212/astro-ph0212571_arXiv.txt": {
        "abstract": "{ We present the results of a 6.4 square degrees imaging survey of the Pleiades cluster in the $I$ and $Z$-bands. The survey extends up to 3 degrees from the cluster center and is $90\\%$ complete down to $I\\simeq 22$. It covers a mass range from 0.03$\\msun$ to 0.48$\\msun$ and yields 40 brown dwarf candidates (BDCs) of which 29 are new. The spatial distribution of BDCs is fitted by a King profile in order to estimate the cluster substellar core radius. The Pleiades mass function is then derived accross the stellar-substellar boundary and we find that, between $0.03\\msun$ and $0.48\\msun$, it is well represented by a single power-law, $dN/dM\\propto M^{-\\alpha}$, with an index $\\alpha=0.60\\pm 0.11$. Over a larger mass domain, however, from 0.03$\\msun$ to 10$\\msun$, the mass function is better fitted by a log-normal function. We estimate that brown dwarfs represent about 25\\% of the cluster population which nevertheless makes up less than 1.5\\% of the cluster mass. The early dynamical evolution of the cluster appears to have had little effect on its present mass distribution at an age of 120~Myr. Comparison between the Pleiades mass function and the Galactic field mass function suggests that apparent differences may be mostly due to unresolved binary systems. ",
        "introduction": "The Initial Mass Function (IMF), i.e. the mass spectrum resulting from complex physical processes at work during star formation, is a formidable constraint for star formation models. Its determination at low stellar and substellar masses is therefore one of the main motivation for the rapidly expanding quest for brown dwarfs. Very low mass stars and brown dwarfs are also prime candidates to investigate the structure and dynamical evolution of large stellar systems, such as e.g. clusters. Searches for the lowest mass, isolated objects have been conducted in the galactic field, in young open clusters and in star forming regions, brown dwarfs being brighter when younger. The identification of substellar candidates is often based on color magnitude diagrams (hereafter CMD) built from deep wide-field imaging surveys. One of the major shortcomings of this selection method is the contamination by older and more massive late-type field dwarfs which may lie in the same region of the CMD.  The Pleiades cluster (RA=$3^{h}46.6^{m}$, Dec=$+24^{o}04'$) is an ideal hunting ground for substellar objects. It is relatively nearby with a distance of about 125 pc, with the apparent magnitude of massive brown dwarfs bright enough for being easily detected on intermediate-size telescopes. Its age of approximately 120 Myr (Stauffer et al. 1998, Martin et al. 1998) makes the lithium test particularly useful for identifying brown dwarfs~: at this age, the lithium depletion boundary (in mass) coincides with the hydrogen burning mass limit (hereafter HBML).  Also, the Pleiades is a rich cluster with about 1200 members and its galactic latitude is relatively high ($b=-23^{o}$) which minimizes the possible confusion between members and red giants from the Galactic disk. Moreover, the cluster motion ($\\mu_{\\alpha}=+19$ mas/yr, $\\mu_{\\delta}=-43$ mas/yr) is large compared to that of field stars so that cluster kinematic studies are a powerful way to recognize true members among the photometric candidates.  \\begin{table*}[htbp] \\centering \\leavevmode \\begin{tabular}{lccccc} \\hline Survey & Area & Completeness & Nb of new & IMF index & Mass range \\\\ & sq. degree & limit & BD candidates & & ($\\msun$) \\\\ \\hline Stauffer et al. (1989) & 0.25 & $I\\sim17.5$ & 4 & & 0.2-0.08 \\\\ Simons \\& Becklin (1992) & 0.06 & $I\\sim19.5$ & 22 &  & 0.15-0.045\\\\ Stauffer et al. (1994) & 0.4 & $I\\sim17.5$ & 2 &  & 0.3-0.075\\\\ Williams et al. (1996) & 0.11 & $I\\sim19$ & 1 &  & 0.25-0.045 \\\\ Festin (1997) & 0.05 & $I\\sim21.6$ & 0 & $\\le1$ &  0.15-0.035 \\\\ Cossburn et al. (1997) & 0.03 & $I\\sim20$ & 1 &  & 0.15-0.04 \\\\ Zapatero et al. (1997) & 0.16 & $I\\sim19.5$ & 9 & $1\\pm0.5$ & 0.4-0.045 \\\\ Festin (1998) & 0.24 & $I\\sim21.4$ & 4 & $\\le1$ &  0.25-0.035 \\\\ Stauffer et al. (1998) & 1. & $I\\sim18.5$ & 3 &  & 0.15-0.035\\\\ Bouvier et al. (1998)$^{a,b}$ & 2.5 & $I\\sim22$ & 13 & $0.6\\pm0.15$ & 0.4-0.04 \\\\ Zapatero-Osorio et al. (1999)$^b$ & 1. & $I\\sim21$ & 41 &  & 0.08-0.035 \\\\ Hambly et al. (1999) & 36. & $I\\sim18.3$ & 6 & $\\le0.7$ &  0.6-0.06 \\\\ Pinfield et al. (2000)$^b$ & 6 & $I\\sim19.6$ & 13 &  & 0.45-0.045 \\\\ Tej et al. (2002) & 7 & $K\\sim15$ & & $0.5\\pm0.2$ & 0.5-0.055 \\\\ Dobbie et al. (2002)$^{b,c}$ & 1.1 & $I\\sim22$ & 10 & 0.8 & 0.6-0.030 \\\\ \\hline \\end{tabular}  \\parbox{0.9\\textwidth}{% $^a$BD candidates have been confirmed by Mart\\'{\\i}n et al. (2000) and Moraux et al. (2001) and the IMF index has been revised to $0.51\\pm0.15$.  $^b$Jameson et al. (2002) have compiled these surveys and using new infrared data they find that the Pleiades mass function is well represented by a power law with index $\\alpha=0.41\\pm0.08$ for $0.3\\msun\\ge M\\ge 0.035\\msun$.  $^c$These authors used the results from Hodgkin \\& Jameson (2000) and Hambly et al. (1999) to conclude that the $\\alpha=0.8$ power law is appropriate from $0.6\\msun$ down to $0.03\\msun$.\\medskip  } \\caption{Previous imaging surveys of the Pleiades conducted for searching brown dwarfs.} \\label{bdsurveys} \\end{table*}  To date, several surveys have been conducted in this cluster. They are summarized in Table~\\ref{bdsurveys} in term of covered area, completeness limit and number of new brown dwarfs candidates. Only surveys with a magnitude limit larger or equal to $I\\simeq17.5$ are listed here (the HBML corresponds to $Ic=17.8$, Stauffer et al. 1998). Starting in 1997\\footnote{Prior to 1997, mass function estimates were based on strongly contaminated samples of brown dwarf candidates.}, several groups reported estimates for the Pleiades mass function in the upper part of the substellar domain, often represented by a single power law $dN/dM \\propto M^{-\\alpha}$ with $0.5\\leq\\alpha\\leq 1.0$ (see Table~\\ref{bdsurveys}).   While these various estimates agree within uncertaintites, the substellar samples on which they rely are still relatively small and, in some cases, not fully corrected for contamination by field dwarfs. Some surveys are also quite limited spatially, and the derived mass function may not be representative of the whole cluster. In order to put the determination of the Pleiades substellar mass function on firmer grounds, we performed a deep and large survey of the Pleiades cluster in the $I$ and $Z$-band, using the CFH12K camera at the Canada-France-Hawaii telescope. The survey covers 6.4 square degrees and reaches up to 3 degrees from the cluster center. It encompasses a magnitude range $I\\simeq 13.5-24$, which corresponds to masses from 0.025$\\msun$ to 0.45$\\msun$ at the distance and age of the Pleiades.   In section 2, we describe the observations and the data reduction. Results in the stellar and substellar domains are presented in Section 3. Forty brown dwarf candidates are identified, 29 of which are new discoveries. We discuss contamination of the photometric samples by field stars in the stellar and substellar domains, investigate the radial distribution of substellar objects, and derive the cluster substellar mass function. In Section 4, we discuss the overall shape of the cluster mass function over the stellar and substellar domains, the possible consequences of early cluster dynamical evolution on its present mass function and the implications for the brown dwarf formation process. A comparison is made between the Pleiades and the Galactic field mass functions, which reveals significant differences at the low mass end, a large part of which is probably attributable to unresolved cluster binaries. ",
        "conclusions": "We have conducted a deep wide field photometric survey of the Pleiades cluster to build a sample of probable cluster members with masses in the range $0.03\\msun$ to $0.48\\msun$.  We have identified 40 brown dwarfs candidates, of which 29 are new discoveries. Taking into account the radial distribution of cluster members, we derive the cluster mass function accross the stellar-substellar boundary. We find that a single power-law $dN/dM \\propto M^{-\\alpha}$ with an index $\\alpha=0.60\\pm0.11$ provides a good match to the cluster mass function in the $0.03-0.48\\msun$ range. This new estimate is based on a survey which combines a large radial coverage of the cluster and a realistic assessment of the contamination by field stars. Furthermore, the survey completely covers the $0.03-0.48\\msun$ mass range, so that the result does not rely on the combination of heterogeneous surveys, as has been the case before. We therefore believe this new estimate is reasonably robust. Small changes may be expected when our survey has been followed up with either infrared photometry and/or proper motions.   Over a larger mass domain, covering almost 3 decades in masses from $0.03\\msun$ to $10\\msun$, we find that the cluster mass function is better fitted by a log-normal distribution with $\\langle M\\rangle\\simeq0.25\\msun$ and $\\sigma_{logM}\\simeq 0.52$. When unresolved Pleiades binaries are taken into account, the log-normal Pleiades mass function is not unlike the Galactic disk mass function. This suggests that the dynamical evolution of the cluster has had yet little effect on its mass content at an age of 120~Myr. It also suggests that the brown dwarf formation process does not lead to the dynamical evaporation of substellar objects at the time the cluster forms."
    },
    "0212/astro-ph0212437_arXiv.txt": {
        "abstract": "We report the first results of the {\\it Chandra} temporal monitoring of the ultra-luminous X-ray sources (ULXs) in the Antennae galaxies (NGC~4038/39). Observations at four different epochs, covering time scales of 2~years to 2~months, show variability in seven out of nine ULXs, confirming that they are likely to be accreting compact X-ray binaries (XRBs). The seven variable ULXs exhibit a variety of temporal and spectral behaviors: one has harder X-ray colors with decreasing luminosity, similar to the black hole binary Cyg X-1, but four other ULXs show the opposite behavior. We suggest that the latter may be black-hole binaries accreting at very high rates. ",
        "introduction": "At a distance of 19~Mpc ($H_o = 75$), NGC~4038/39 have long been studied as the nearest examples of a galaxy pair undergoing a major merger. Recently, high resolution {\\it HST} (e.g., Whitmore et al.\\ 1999), {\\it Chandra} (Fabbiano et al.\\ 2001), and {\\it VLA} (Neff \\& Ulvestad 2000) observations have produced exquisitely detailed images of these galaxies, providing an in-depth probe of their young stellar population. In the X-ray band (0.1--10~keV), the first {\\it Chandra} observation of this system in December 1999 (Fabbiano et al.\\ 2001) revealed a population of  extraordinarily luminous point-like sources. Nine of these sources, which have hard spectra typical of X-ray binaries (XRBs), were detected with luminosities $L_X > 10^{39} \\rm ergs~s^{-1}$ $(H_o=75)$, significantly exceeding the Eddington luminosity of a spherically accreting neutron star (Zezas et al.\\ 2002a). This type of Ultra Luminous X-ray source (ULX) has been reported in other nearby galaxies (e.g., Makishima et al.\\ 2000), but never in such copious numbers. ULXs have been alternatively explained with massive black hole binaries ($M > 10 - 100 M_{\\odot}$; e.g. Fabbiano 1989; Makishima et al.\\ 2000), young Supernova remnants (e.g. Fabian \\& Terlevich 1996), or beamed XRBs (King et al.\\ 2001). The ULXs of The Antennae appear significantly displaced from nearby young stellar clusters (Zezas et al.\\ 2002b), suggesting that a significant fraction may not have a very massive counterpart (Zezas \\& Fabbiano 2002).  Given their large number of ULXs and our thorough knowledge of their stellar populations, The Antennae provide an exceptional environment for studying ULXs and comparing them to the more normal XRB population. To this end, a large {\\it Chandra} monitoring campaign of The Antennae is under way. Here we report our first results, by comparing three recent observations with the original December 1999 data. ",
        "conclusions": "We have found widespread variability of all non-nuclear ULXs in the Antennae galaxies. While one of these sources follows a Cyg~X-1 like high/soft--low/hard behavior, four other sources instead become softer with decreasing luminosity. We suggest that black hole binaries with very high accretion rates could explain the low/soft correlation. In all cases beaming due to thick warped disks could boost the luminosity to super-Eddington values."
    },
    "0212/hep-ph0212387_arXiv.txt": {
        "abstract": "Selected topics in  neutrino astrophysics are reviewed. These include  the production of low energy neutrino flux from cores of collapsing stars and the expected high energy neutrino flux from some other astrophysical sites such as the galactic plane as well as the center of some distant galaxies. The expected changes in these neutrino fluxes because of neutrino oscillations during their propagation to us are described. Observational signatures for these neutrino fluxes with and without neutrino oscillations are discussed. ",
        "introduction": "In the standard model of the electro weak interactions, the lepton masses and the values of other parameters such as weak mixing angle, couplings, etc. are arbitrary and are therefore determined by experiments. These parameters are independent of each other and can not be determined uniquely, while the neutrino is taken to be massless because of maximal parity violation. The masslessness of the neutrino however does not follow from any other theoretical ground unlike the local gauge invariance for the photon \\cite{Fukugita:wx,Kim:dy,djouadi}.   In order to search for physics {\\tt beyond} the standard model of particle physics, light neutrino masses are incorporated in the extensions of the standard model. Then, on quite general grounds one expects that neutrinos will also possess non-zero magnetic moments. Experimentally, one places finite non-vanishing upper bounds on measurable neutrino masses and magnetic moments. Presently, there is some indirect experimental evidence for the masslessness of the neutrino \\cite{Valle:2002tm}.   Massive neutrinos quite likely mix. The mixing of quarks is an established fact and because of quark-lepton symmetry, it is natural to assume that leptons exhibit mixing as well. An additional argument, in this sense, is provided by grand unification models, in which quarks and leptons are described in a unified manner.  Mixing means that $\\nu_{e}$, $\\nu_{\\mu}$ and $\\nu_{\\tau}$, i.e., the states created in weak interactions, are different from the states $\\nu_{1}$, $\\nu_{2}$ and $\\nu_{3}$ that have definite masses. The neutrinos $\\nu_{e}$, $\\nu_{\\mu}$ and $\\nu_{\\tau}$ are orthogonal combinations of $\\nu_{1}$, $\\nu_{2}$ and $\\nu_{3}$ with different phases between them. In addition, the sterile neutrinos may mix with these neutrinos.  Neutrino mixing has, as its consequence, neutrino oscillations, i.e., the process of periodic (complete or partial) conversion of neutrinos of one type into another, for instance, $\\nu_{e}\\to \\nu_{\\mu}\\to \\nu_{e}\\to ... $ . The components $\\nu_{i}$ of a mixed neutrino have different masses and hence different phase velocities. It follows that the phase difference caused by the mass difference between the $\\nu_{i}$ vary monotonically during the propagation. This phase change manifests itself as neutrino oscillations.   If neutrino oscillations do occur in vacuum, matter can enhance their depth (probability amplitude) up to a maximal value \\cite{Mikheev:qk}. That is, a monotonic change of density may lead to resonant conversions between various neutrino flavors. This follows from the fact that when neutrinos propagate through a monotonically changing density medium, $\\nu_{e}$ and $\\nu_{\\mu}$ ($\\nu_{\\tau}$) feel different potentials, because $\\nu_{e}$ scatters off electrons via both neutral and charged currents, whereas $\\nu_{\\mu}(\\nu_{\\tau})$ scatters off electrons only via the neutral current. This induces a coherent effect in which maximal conversion of $\\nu_{e}$ into $\\nu_{\\mu}$ take place (even for a rather small intrinsic mixing angle in the vacuum), when the phase difference arising from the potential difference between the two neutrinos cancel the phase caused by the mass difference in the vacuum \\cite{Wolfenstein:1977ue}.   During nearly past half a century, the empirical search for neutrinos has spanned roughly four orders of magnitude in neutrino energy $E$, from $\\sim $ MeV up to $\\sim  10^{6}$ MeV. The lower energy edge corresponds to the Solar neutrinos, whereas the upper energy edge corresponds to the Atmospheric neutrinos. A detailed early description of the Solar neutrino search can be found in \\cite{Bahcall:gw}, whereas for recent status, see \\cite{lipari}. The aspects of neutrino production in Atmosphere of earth related to neutrino oscillation studies are recently reviewed in \\cite{Gaisser:2002jj}. The intermediate energy range corresponds to terrestrial neutrinos such as from reactors (and accelerators) and the supernova neutrinos. Thus, obviously either going in energy range below these values or {\\tt above} are the available frontiers.   For a general introduction of the possibility of having neutrinos with energy $> 10^{6}$ MeV, see \\cite{Bahcall:iz}, the upper energy edge for these high energy neutrinos is limited only by the concerned experiments. More detailed general discussions in the context of high energy neutrinos can be found in  \\cite{Ginzburg:sk,Learned:sw}. Despite the availability of the somewhat detailed discussion of progress cited in the last reference, the field of high energy neutrino astrophysics is still passing through its initial stage of development.   The above mentioned empirical search has already given us quite useful insight into neutrino intrinsic properties such as mass and mixing. Massive neutrinos and their associated properties such as Dirac or Majorana character of their mass, their mixings and magnetic moments can have important consequences in astrophysics. In these two lectures, I elaborate some of the selected  consequences and the constraints implied by these consequences on neutrino properties as well as the insight that one may gain about the nature of the astrophysical or/and cosmological sites and the interactions that produce these neutrinos. The explanation of observed Solar $\\nu_{e}$ deficit relative to its production value in the core of the Sun, via $\\nu_{e}\\to \\nu_{\\mu }, \\nu_{\\tau} $ conversion is an interesting example in this context \\cite{Smirnov:2002in}.   The general plan of the lectures is as follows. In the first lecture, I elaborate these aspects for the low energy neutrino flux ($E \\sim $ MeV) emitted during the gravitational collapse of stars which may be accompanied by supernovae phenomenon. In the second lecture, I elaborate these for the expected high energy neutrino flux ($E \\geq 10^{6}$ MeV) from some representative examples of remaining cosmos around us such as our galactic plane. This includes the one arising from the interaction of ultra high energy cosmic ray flux ($E \\geq 10^{12}$ MeV) with the matter and radiation inside the sources (such as center of our and other galaxies) of as well as during propagation of ultra high energy cosmic ray flux to us.  \\pagebreak ",
        "conclusions": "Detailed study of neutrino fluxes from different astrophysical sites such as the cores of collapsing stars, the galactic plane, as well as other more far away anticipated astrophysical and cosmological sites will provide valuable information about the neutrino intrinsic properties and the site itself.   In collapsing stars, such studies can constrain the role of supernova magnetic field in mixed neutrino propagation. In particular, if a neutrino spin flavor conversion occurs in the isotopically neutral region in a collapsing star, the resulting changes in neutrino flux spectrum are sensitive to rather small values of neutrino magnetic moment, $\\mu $, such as $\\mu \\leq 10^{-13}\\mu_{B}$. In this context, in the first lecture, after a brief introduction of the main neutrino production mechanisms during the supernova stage, the description for neutrino spin flavor conversions is elaborated under the assumption of small vacuum mixing. This includes a discussion that incorporates the possibility of violation of adiabaticity in neutrino spin flavor conversions. Results of its implications for supernova neutrinos are summarized.    For other more distant and energetic sources, the prospective high energy neutrino observation will provide clues for a solution of the long standing problem of origin of observed ultra high energy cosmic rays. The (non) observation of high energy neutrinos will also help to better model the underlying physics of the far away astrophysical and cosmological  sites. In this context, in the second lecture, the present motivations for their searches are presented, with a description of their possible connection with ultra high energy cosmic rays. Main high energy neutrino production mechanisms are summarized via a simple classification flow chart. Three flavor neutrino oscillation description is reviewed and its implications for the three relative ratios of high energy neutrinos are given. Furthermore, the essentials of prospective observations of high energy neutrinos are briefly described including the possible relevant observational signatures.   The author thanks Physics Division of NCTS for financial support and G.-L. Lin and T.-W. Yeh for comments.  \\pagebreak"
    },
    "0212/astro-ph0212167_arXiv.txt": {
        "abstract": "{We analyse the \\chandra dataset of the galaxy cluster MS1008.1$-$1224 to recover an estimate of the gravitating mass as function of the radius and compare these results with the weak lensing reconstruction of the mass distribution obtained from deep FORS1-VLT multicolor imaging.  Even though the X-ray morphology is disturbed with a significant excess in the northern direction suggesting that the cluster is not in a relaxed state, we are able to match the two mass profiles both in absolute value and in shape within $1 \\sigma$ uncertainty and up to 1100 $h_{50}^{-1}$ kpc. The recovered X-ray mass estimate does not change by using either the azimuthally averaged gas density and temperature profiles or the results obtained in the northern sector alone where the signal-to-noise ratio is higher. ",
        "introduction": "As the largest virialized objects in the Universe, galaxy clusters are a powerful cosmological tool once their mass distribution is univocally determined.  In the recent past, there have been several claims that cluster masses obtained from X-ray analyses of the intracluster plasma, taken to be in hydrostatic equilibrium with the gravitational potential well, are significantly smaller (up to a factor of two; but see Wu et al.\\ 1998; Allen 1998; B\\\"ohringer et al.\\ 1998; Allen et al.\\ 2001) than the ones derived from gravitational lensing (see Mellier 1999 for a review).  \\begin{figure} \\begin{center} \\epsfig{file=Ek201_f1.ps,width=0.45\\textwidth} \\caption{An adaptive smoothed exposure-corrected image in the Chandra [0.5--2] keV band is overplotted to the R-band image (the white areas in the upper end show masked regions). The contours are spaced with steps of $\\log 2$ from the minimum of $2.1 \\times 10^{-9} \\mbox{ photon sec}^{-1} \\mbox{ cm}^{-2} \\mbox{ arcsec}^{-2}$.} \\end{center} \\label{opt_xray} \\end{figure}  In this paper we report on the mass distribution of the cluster MS1008.1$-$1224 by combining the results from weak lensing analysis of deep FORS1-VLT images with those obtained from a spatially-resolved spectroscopic X-ray analysis of a \\chandra observation. MS1008.1$-$1224 is a rich galaxy cluster at redshift 0.302 that has been part of the Einstein Medium Sensitivity Survey sample (Gioia \\& Luppino 1994) and of the CNOC survey (Carlberg et al.\\ 1996).  Lombardi et al.\\ (2000) presented a detailed weak lensing analysis of the FORS1-VLT data.  Figure~\\ref{opt_xray} shows the X-ray isophotes overplotted to the optical V-band image. In the following we adopt the conversion $1 \\mbox{ arcmin} = 376 \\mbox{ kpc}$ ($z=0.302$, $H_0 = 50 \\, h_{50} \\mbox{ km s}^{-1} \\mbox{ Mpc}^{-1}$, $\\Omega_{\\rm m} = 1 - \\Omega_{\\Lambda} = 0.3$) and quote all the errors at $1 \\sigma$ ($68.3\\%$ confidence level). ",
        "conclusions": "The differential X-ray best-fit mass model has been weighted by the relative portion of the shell observed in each ring and, then, cumulated up to the radius of 1020 kpc (the outer radius of our spatially resolved spectroscopy). In Fig.~\\ref{fig:mass}, we plot and compare the projected mass profiles of the galaxy cluster MS1008.1$-$1224 obtained from both the spatially resolved spectral analysis of the \\chandra observation and the weak lensing analysis of FORS1-VLT multicolor imaging. The two independently reconstructed mass profiles agree very well within $1 \\sigma$ uncertainty both in absolute values and in the overall shape of the profile. Note that the mass center is fixed to the peak of the lensing map density that is consistent with the X-ray peak as discussed in Sect.~2. This result does not change when we consider the different density and temperature profiles observed in the northern region where a significant surface brightness excess is located and a higher signal-to-noise ratio is available, arguing for the robustness of the X-ray mass estimates once gas density and temperature distributions can be properly mapped."
    },
    "0212/astro-ph0212484_arXiv.txt": {
        "abstract": "The spectra of stars with the B[e] phenomenon are dominated by features that are related to physical conditions of circumstellar material around these objects and are not intrinsic to the stars. Previous studies have shown that emission lines present in the optical spectra of these objects are formed in an equatorial rotating disk. The analysis of high and low resolution spectra, obtained by us with 1.52 telescope in ESO for some Southern Galactic B[e] Stars, can give us information about the structure and velocity of the disk. We will describe the analysis of the unclassified B[e] star Hen 2-90. ",
        "introduction": "From a sample of stars with the B[e] phenomenon (Lamers et al. 1998) observed by us with the FEROS and B\\&C spectrograph at 1.52 telescope in ESO (La Silla, Chile - agreement ESO/ON), we have analyzed the circumstellar medium around {\\index Hen 2-90}, a unclassified B[e] star (that is sometimes classified in the literature as a compact planetary nebulae) considering the presence of a rotating disk. Its presence is suggested by images taken with WFPC2 in HST (Sahai \\& Nyman 2000) and also by the presence of double peaked profiles in its spectrum. Our analysis is based on the comparison between the observed luminosity of [S\\,{\\sc ii}] lines and H$\\alpha$ with those predicted by a model that will be described below. ",
        "conclusions": "We have found a $\\dot{M}$ $\\sim$ 10$^{-9}$ M$_\\odot$ year$^{-1}$ for the disk region, if all S is S\\,{\\sc ii}. This value is low compared with other compact planetary nebula. If the S\\,{\\sc ii} fraction (S\\,{\\sc ii}/S) is only 10$^{-3}$, then the $\\dot{M}$ is $\\sim$ 10$^{-7}$ M$_\\odot$ year$^{-1}$, because we have found that $\\dot{M}$ is proportional to (S\\,{\\sc ii}/S)$^{-0.67}$ (Borges Fernandes et al. 2003).  We will apply this model for other stars in our sample. For this, we will improve the model including other elements like O and N."
    },
    "0212/astro-ph0212351_arXiv.txt": {
        "abstract": "We present a proper motion survey over about 200 square degrees towards the NGP, based on the material used for the construction of the GSC-II, that we are using to study the vertical structure and kinematics of the Galaxy.  In particular, we measured the rotation velocity of the halo up to 10 kpc above the galactic plane traced by a sample of RR Lyr{\\ae} and BHB giants for which radial velocities were used to recover the complete distribution of the spatial velocities.  Finally, the impact of astrometric and spectroscopic GAIA observation are discussed. ",
        "introduction": "It is generally accepted that the Galaxy is constituted by four discrete main components, the {\\it bulge}, the {\\it thin disk}, the {\\it thick disk} and the {\\it halo}, which are characterized by distinctive stellar populations in terms of spatial distribution, kinematics properties, metallicity and age. A detailed knowledge of such galactic components is essential to achieve a complete description of the Milky Way, as well as of the various processes and evolutionary phases which occurred during the history of our Galaxy (and other galaxies too) and that are responsible for the existence and the properties of those components we observe today.   Ground based surveys providing photometry, proper motions and/or spectroscopic observations have been carried out in the past to study the structure and kinematics of Galactic populations, and these will continue and be extended in the next years thanks to availability of all-sky catalogs such as GSC-II and USNO-B, based on multi-epoch photographic surveys, as well as other current and future photometric and spectroscopic surveys (eg.\\ 2MASS, DENIS, SDSS, EIS, HES, HK, VST, etc.) that  benefit from the availability of dedicated scanning and large field cameras % as well as  large multi-fiber spectrographs (eg. 2dF, 6dF, FLAMES).  Clearly, the GAIA mission will provide an enormous contribution to the understanding of the formation and evolution of the Galaxy, thanks to very accurate astrometric parameters complemented by photometric and spectroscopic measurements which will permit the direct determination of the 6D phase space distribution and the chemical abundance of large and complete samples of stellar tracers belonging to the different galactic components.  Here we describe a new project which combines astrometric, photometric and spectroscopic data in order to investigate the kinematics of the outer halo by means of velocities derived from proper motions and radial velocities for a sample of RR-Lyr{\\ae} and BHB giants towards the NGP. ",
        "conclusions": "Large field proper motion surveys ($\\sigma_\\mu\\approx$ 1-10 mas/yr) based on photographic surveys digitized with fast measuring machines (APS, MAMA, PDS, PMM, SuperCOSMOS) are still useful tools for the study of the structure and kinematics of the galactic stellar populations, especially when combined with radial velocities and chemical abundances from spectro-photometric observations. In particular, we have shown how bright halo tracers, such as BHB giants and RR Lyr{\\ae} stars with $V\\la 16$-17, can be used to investigate {\\it in situ} the halo kinematics up to $z\\approx 5$-6 kpc via photometric parallaxes if proper motions with $\\sigma_\\mu \\sim 1$-2 mas/yr accuracy and radial velocities with precision of about $\\sigma_{\\rm Vr}\\sim 10$-20 km~s$^{-1}$ are available.  For statistical analysis of large samples, systematic errors on the proper motions (and $M_V$ too) are critical and can result in biased mean velocities and dispersions. These problems will be dramatically reduced by the astrometric parameters and spectroscopic $V_r$  provided by GAIA, which will permit to determine $(a)$ direct distances from parallaxes, $(b)$ radial and tangential velocities with an accuracy of a few tens of km~s$^{-1}$ to $d\\sim 10$-15 kpc for $(c)$ unbiased and larger sets of tracers identified by means of its spectro-photometric and astrometric data."
    },
    "0212/astro-ph0212398_arXiv.txt": {
        "abstract": "Recent numerical simulations of magnetic reconnection in two dimensions have shown that, when the resistivity is strongly localized, the reconnection region develops a Petschek-like structure, with the width of the inner diffusion region being of the order of the resistivity localization scale. In this paper, we combine this fact with a realistic model for locally-enhanced anomalous resistivity generated by current-driven microturbulence. The result is a qualitative model of the reconnection layer where the size of Petschek's diffusion region and, therefore, the final reconnection rate are determined self-consistently in terms of the main parameters of the functional dependence of anomalous resistivity on the electric current density. We then consider anomalous resistivity due to ion-acoustic turbulence as a particular case. This enables us to express the reconnection region's parameters directly in terms of the basic parameters of the plasma. Finally, we apply this reconnection model to solar flares and obtain specific predictions for typical reconnection times, which are very consistent with observations. ",
        "introduction": "\\label{sec-intro}   Magnetic reconnection is a basic plasma physics phenomenon of tremendous importance in many astrophysical systems (Tsuneta~1996; Kulsrud~1998), as well as in some laboratory plasma devices (Yamada~et~al. 1997), including tokamaks (Kadomtsev~1975; Yamada~et~al. 1994). It has been studied extensively over the past five decades (e.g., Giovannelli~1946; Vasyliunas~1975; Biskamp~2000). Historically, the first (and perhaps still the most important) application of magnetic reconnection has been to explain the solar flare phenomenon, and it is in this context that the earliest reconnection models have been developed.  The first theoretical model of magnetic reconnection was developed by Sweet (1958) and by Parker (1957, 1963). In this model, magnetic field is frozen into the plasma everywhere except in a very thin layer where the current density is so high that resistive effects become important no matter how small the resistivity is. It is inside this thin current layer that the actual breaking (and reconnecting) of the magnetic lines of force takes place, accompanied by a violent release of enormous amounts of magnetically stored energy, thus leading to the observed flare. The Sweet--Parker theory predicts that the current layer thickness, $\\delta_{\\rm SP}$, scale as $\\delta_{\\rm SP} \\sim L/\\sqrt{S}$, where $L$ is the global system size (typically of order $10^9$ cm in the solar corona) and $S\\equiv L V_A/\\eta$ is the global Lundquist number (here $V_A$ is the Alfv{\\'e}n velocity and~$\\eta$ is the magnetic diffusivity; in the rest of this paper we shall refer to $\\eta$ as the resistivity). Correspondingly, the typical reconnection timescale is found to be $\\tau_{\\rm rec}\\sim \\tau_A(L) \\sqrt{S}$, where $\\tau_A(L)\\equiv L/V_A$ is the global Alfv{\\'e}n transit time. It has been immediately realized that the resulting reconnection time turns out to be too long; in the solar corona one typically has $S\\sim 10^{12}-10^{14}$ and $\\tau_A(L)\\sim 1$~sec, which leads to $\\tau_{\\rm rec}$ of the order of a few months. This is in sharp contrast with the typical observed solar flare duration of order $10^2 - 10^3$~sec. Thus, since its early years, the main thrust of magnetic reconnection research has been to explain reconnection rates that are much faster than the Sweet--Parker theory predicts.  Starting from the 1960s, two major routes toward faster reconnection were proposed. One of them was to use the so-called {\\it anomalous resistivity} instead of the classical Spitzer resistivity used in the original Sweet--Parker model (Coppi \\& Friedland 1971; Smith \\& Priest 1972; Coroniti \\& Eviatar 1977; Kulsrud~1998). The idea was that, as a current layer forms, its thickness becomes so small, and hence the current density becomes so high, that the drift velocity of the current-carrying electrons exceeds a certain threshold, such as the ion-sound or the electron thermal speed. This leads to the excitation of current-driven kinetic microturbulence, which, in turn, provides a more efficient (compared to particle-particle collisions) mechanism for the scattering of electrons (via wave-particle interactions). As a result, one ends up with a greatly enhanced effective resistivity and, hence, a greatly reduced effective Lundquist number. When substituted into the Sweet--Parker scaling, this results in a greatly enhanced reconnection rate. In fact, controlled laboratory studies of magnetic reconnection have shown a good agreement with a simple Sweet--Parker model augmented with some (experimentally measured) anomalously enhanced resistivity (Ji~et~al. 1998, 1999). On the other hand, in the solar flare context, the Sweet--Parker model with anomalous resistivity gives typical reconnection times of the order of several hours (e.g., Kulsrud~1998), a great improvement over Spitzer resistivity. Among various anomalous resistivity mechanisms, the one most frequently quoted has been the {\\it ion-acoustic turbulence} (IAT). Rapid magnetic energy dissipation due to the IAT-driven anomalous resistivity has in fact also been invoked to explain coronal heating (Rosner~et~al.~1978).  Another possibility leading to shorter reconnection times was proposed by Petschek (1964). His very elegant model makes use of a somewhat more complicated reconnection layer geometry, while still relying on simple resistive magnetohydrodynamics (MHD) without invoking any new physics at the microscopic level. The {\\it Petschek model} actually does not predict a unique configuration of the reconnection layer and a unique reconnection rate. Instead, this model encompasses an entire one-parametric family of solutions. Each of these configurations has at its center a small Sweet--Parker-like layer, called the {\\it diffusion region}, and four standing slow-mode shocks emanating from the ends of this central layer. The members of this family of solutions can be labeled by the width (or length) $\\Delta$ of the inner diffusion region. As Petschek noticed, in the Sweet--Parker model the reconnection process had been slowed down by the very large aspect ratio $L/\\delta$ of the reconnection layer. He suggested that the width~$\\Delta$ of the layer's diffusion region does not have to be as large as the global size~$L$; if $\\Delta$ can be made sufficiently short, the reconnection process will go much faster than in the Sweet--Parker model. The maximum value of~$\\Delta$ corresponds to the Sweet--Parker solution, which is, therefore, just one of the family members. The reconnection rate ranges from the slowest (Sweet--Parker) rate to the so-called maximum Petschek rate, which scales as $1/\\log{S}$. We thus see that the relatively strong square-root dependence on the resistivity, characteristic for the Sweet--Parker model, is replaced here by a much weaker logarithmic dependence; even for a very large $S\\sim 10^{14}$, the resulting reconnection timescale turns out to be reasonably short, of order $10^2\\, \\tau_A(L)$.  This model had remained the favorite model of reconnection until 1980s, when two-dimensional (2D) resistive-MHD numerical simulations by Biskamp (1986) showed that, in the case of a spatially uniform resistivity, a Petschek-like configuration fails to form and that a long (of order~$L$) current layer tends to form instead, consistent with the Sweet--Parker picture. This finding has been confirmed in a number of numerical simulations performed by several other groups (Scholer~1989; Ugai 1992, 1999; Yokoyama \\& Shibata 1994; Uzdensky \\& Kulsrud~2000; Erkaev~et~al. 2000, 2001). A theoretical explanation has been put forward by Kulsrud (2001). He noticed that, when the resistivity is uniform, the relatively large transverse magnetic field that is needed to support the standing shocks in the Petschek model is rapidly swept out of the diffusion region by the downstream flow, while its regeneration due to the nonuniform merging is not fast enough. As a result, the diffusion region's size~$\\Delta$ (which Kulsrud calls~$L'$) increases until it reaches the global scale~$L$ and the reconnection rate, correspondingly, slows down to the usual Sweet--Parker rate. This explanation has been confirmed numerically by Uzdensky \\& Kulsrud (2000) (see also Kulsrud~1998).  It has been noticed, however, that the key assumption leading to the above conclusion was the assumption of uniform resistivity (which is a very common assumption in numerical simulations in general). Simulations featuring a {\\it non-uniform resistivity} have shown that a Petschek-like structure does form and can be stable whenever the resistivity is locally enhanced in some small region near the X-point at the center of the reconnection region (Ugai \\& Tsuda 1977; Sato \\& Hayashi 1979; Ugai 1986, 1992, 1999; Scholer~1989; Yokoyama \\& Shibata 1994; Erkaev et al. 2000, 2001; Ugai \\& Kondoh 2001; Biskamp \\& Schwarz 2001). A nice plausible theoretical explanation of this phenomenon has again been provided by Kulsrud (2001), who analyzed how a locally enhanced resistivity may lead to a more efficient regeneration of the transverse magnetic field through non-uniform merging and thus to the sustainment of Petschek's shocks. In addition, recent numerical work by Erkaev et~al. (2000, 2001) and by Biskamp \\& Schwarz (2001) has shown that the particular Petschek-like configuration that forms in the localized-resistivity situation is characterized by the width~$\\Delta$ of the inner diffusion region being of the order of the resistivity localization scale, which in this paper we shall call~$l_\\eta$ (Erkaev~et~al. 2000, 2001; Biskamp \\& Schwarz 2001). We thus see that resistivity-localization mechanism plays a crucial role in determining the final reconnection rate.  From the point of view of physical reality, the main motivating force behind these localized-resistivity studies has been the idea that anomalous resistivity, being such a sensitive function of the local current density, may in fact be triggered only in a small neighborhood of the X-point, where the current density is highest. In other words, anomalous resistivity can enhance reconnection rate not only directly (by simply being higher than the collisional resistivity), but also indirectly, via enabling the Petschek mechanism (by being strongly localized). This is one of the key ideas behind the so-called spontaneous fast reconnection model suggested by Ugai (1986, 1992, 1999; also see Ugai \\& Tsuda 1977; Yokoyama \\& Shibata 1994; Ugai \\& Kondoh 2001) on the basis of numerical simulations.  Up until now, however, with the notable exception of the work by Kulsrud (2001), there have been no analytical attempts to combine the Petschek reconnection model with any physically realistic model of anomalous resistivity, which would provide unique theoretical predictions for the main parameters of the reconnection region, such as the width of the diffusion region and the reconnection rate, in terms of the basic parameters of the plasma. The goal of this paper is to remedy this situation by attempting to build a simple theoretical framework incorporating a particular anomalous resistivity mechanism (ion-acoustic turbulence).  In \\S~\\ref{sec-model} we present our model of the Petschek reconnection layer with a generic form of anomalous resistivity. In particular, in \\S~\\ref{subsec-ingredients} we describe the basic elements of the model, e.g., the Sweet--Parker relationships for the inner diffusion region and the functional dependence of anomalous resistivity, $\\eta$, on the current density~$j$. In \\S~\\ref{subsec-stability} we discuss possible solutions and explore their stability. In \\S~\\ref{subsec-rate} we calculate the reconnection rate in terms of the model parameters. \\S~\\ref{sec-IAT} is devoted to the specific case of anomalous resistivity due to the IAT; in this section we make use of the well-developed theory of ion-acoustic turbulence to express all major reconnection layer parameters, including the reconnection rate, in terms of the basic plasma parameters (such as the magnetic field strength, plasma density, and the electron and ion temperatures). We apply the obtained results to solar flare environment in~\\S~\\ref{sec-flares} and find a very reasonable agreement in terms of general timescales. We list our conclusions in~\\S~\\ref{sec-conclusions}. ",
        "conclusions": "\\label{sec-conclusions}  In this paper we have presented a model of magnetic reconnection in the presence of a current-driven enhanced anomalous resistivity. This is a very simplistic, crude model that aims at predicting the qualitative behavior of the system and the scaling of the reconnection rate with various plasma parameters, while treating numerical factors of order one only very approximately.  In this model we have combined the following three ingredients. The first one is the observation, derived from several recent resistive-MHD numerical simulations (Erkaev~et~al. 2000, 2001; Biskamp \\& Schwarz 2001), that whenever the resistivity is strongly localized, the reconnecting system will develop a Petschek-like configuration, with the width of the inner diffusion region of the order of the resistivity localization scale. The second ingredient of our model is the Sweet--Parker model (Sweet~1958; Parker 1957, 1963) for the diffusion region of a Petschek configuration (Petschek~1964). Finally, the third ingredient is a physically realistic model for a current-driven anomalous resistivity expressed as a function~$\\eta(j)$, which exhibits two characteristic features. The first feature is a sudden jump of~$\\eta$ from a small collisional value~$\\eta_0$ to a much larger value~$\\eta_1$ as soon as~$j$ exceeds a threshold~$j_c$. The second feature is a subsequent linear growth $\\eta\\propto j$ for~$j>j_c$. This choice is motivated by the theory of anomalous resistivity due to ion-acoustic turbulence, which has been developed in detail over the last 40 years (see, e.g., Bychenkov et al. 1988).  Note that the anomalous resistivity function adopted in this paper becomes very sensitive to electric current density when the latter exceeds some threshold value~$j_c$; this makes it possible for the resistivity to be enhanced only in a small region, which, in turn, leads to the development of a Petschek-like configuration. Thus, our model is characterized by a reconnection rate that is enhanced (with respect to the classical, collisional-resistivity Sweet--Parker rate) by a combined action of anomalous resistivity and of the Petschek mechanism. It is important to realize that the role of anomalous resistivity in the acceleration of the reconnection process is two-fold: in addition to its direct action (lowering the global Lundquist number $S\\equiv V_A L/\\eta$), it accelerates reconnection indirectly, by turning on the Petschek mechanism.  The width of the inner diffusion region of the Petschek model, and thus the resistivity localization scale, are determined self-consistently when all the ingredients of the model are taken into account.  Based on our stability analysis of two possible Petschek-like states, we predict that the system will evolve towards a certain stable Petschek-like configuration. This stable configuration is characterized by the central current density $j_0$ and the central resistivity $\\eta(j_0)$ exceeding $j_c$ and $\\eta_1$, respectively, by a finite factor of order one. The reconnection velocity then scales as $V_{\\rm rec} \\sim V_A/S_*=\\eta_1/\\delta_c$, where $\\delta_c=cB_0/4\\pi j_c$ is the critical thickness of the layer.  We then consider (in \\S~\\ref{sec-IAT}) the case of anomalous resistivity due to ion-acoustic turbulence, as an important specific example. We derive very simple expressions for the parameters of the reconnection system (e.g., the width~$\\Delta$ of the diffusion region, reconnection velocity $V_{\\rm rec}$, and the reconnection time scale $\\tau_{\\rm rec}$) in terms of the basic plasma parameters $n_e$, $T_e$, $T_i$, and~$B_0$.  Finally, in \\S~\\ref{sec-flares}, we apply our model to the solar flare environment. We note that reconnection process will lead to significant electron heating, so that the electron pressure at the center of the reconnection layer may become comparable to the pressure of the reconnecting magnetic field outside the layer. Based on our model, we obtain typical reconnection times of order $10^2-10^3$ sec; this is short enough to explain the very fast time scale of impulsive flares. We note however that, as a result of the plasma heating inside the reconnection layer, the thickness $\\delta$ of the diffusion region quickly becomes comparable to, or even smaller than, the ion skin-depth, $d_i \\equiv c/\\omega_{pi}$. At these scales, new physical processes, described by Hall MHD, may come into play (Drake~et~al. 1994; Biskamp~1997; Bhattacharjee~et~al. 2001) even before the IAT develops and anomalous resistivity becomes important.   I am grateful to S.~Boldyrev, H.~Li, R.~Kulsrud, and R.~Rosner for some very helpful discussions and interesting comments. I would like to acknowledge the support by the NSF grant NSF-PHY99-07949."
    },
    "0212/astro-ph0212021_arXiv.txt": {
        "abstract": "In direct imaging, broad band flat-fields can easily be applied to correct deviation from uniform sensitivity across the detector field. However for slitless spectroscopy data the flat field is both field and wavelength dependent. The effect of the wavelength dependent flat field for the slitless G800L grism mode of the ACS Wide Field Channel (WFC) has been investigated from observations of a flux calibrator at different positions in the field. The results of various flat-fielding schemes are presented including application of a flat field cube derived from in-orbit broad band filter flat fields. Excellent results are reported with deviations in the extracted spectra $<2\\%$ across the WFC field. ",
        "introduction": "New ACS in-orbit filter flat-fields, were recently constructed using globular cluster observations (Mack et al., 2002).  These flat-fields show that there is a significant amount of large scale structure which varies from one broad band filter to another (i.e with wavelength). Flat-fielding ACS G800L grism data is therefore required in order to be able to derive a unique G800L sensitivity curve which is the same at all positions on the detector. We have used these flat-fields and have fitted them using a 3rd order polynomial as a function of wavelength and at each pixel. The result of this fit was the creation of a data cube enabling aXe, the ST-ECF Slitless Extraction Software (Pirzkal et al., 2001), to compute the proper flat-fielding coefficient at any pixel position and wavelength.  Unfortunately, the in-orbit LP flats do not offer an ideal sampling of the wavelength dependence of the flat-field because: 1) they are broad; 2) only a limited number of filters are available; and, 3) there is no filter which reaches wavelengths beyond 8500\\AA~ (while the G800L grism mode remains sensitive to about 10,000\\AA). ",
        "conclusions": "We have successfully constructed a G800L flat-field data cube which allows one to reach flux calibration accuracies of about $1\\%$ at wavelengths ranging from 6000\\AA~ to 9000\\AA~ and across most of the WFC field of view . This modified G800L-corrected flat-field cube can be used directly by the extraction software aXe to extract un-drizzled, non-geometrically corrected grism observations. It will be made available from the ST-ECF ACS spectroscopy group at http://www.stecf.org/instrument/acs/ ."
    },
    "0212/gr-qc0212066_arXiv.txt": {
        "abstract": "New methods are presented which enables one to analyze the thermodynamics of systems with long-range interactions.  Generically, such systems have entropies which are non-extensive, (do not scale with the size of the system).  We show how to calculate the degree of non-extensivity for such a system.  We find that a system interacting with a heat reservoir is in a probability distribution of canonical ensembles. The system still possesses a parameter akin to a global temperature, which is constant throughout the substance. There is also a useful quantity which acts like a {\\it local temperatures} and it varies throughout the substance.  These quantities are closely related to counterparts found in general relativity. A lattice model with long-range spin-spin coupling is studied. This is compared with systems such as those encountered in general relativity, and gravitating systems with Newtonian-type interactions. A long-range lattice model is presented which can be seen as a black-hole analog. One finds that the analog's temperature and entropy have many properties which are found in black-holes. Finally, the entropy scaling behavior of a gravitating perfect fluid of constant density is calculated. For weak interactions, the entropy scales like the volume of the system.  As the interactions become stronger, the entropy becomes higher near the surface of the system, and becomes more area-scaling. ",
        "introduction": "In the study of thermodynamics, it is almost always implicitly assumed that the system does not possess long-range interactions.  Very little is known about the thermodynamics of systems which do possess long-range interactions, except in special cases such as plasmas where the electromagnetic interactions are screened, or systems which have no overall charge\\cite{landlieb}.  In both these instances, one can use standard thermodynamics, since effectively, there is no long-range interaction. If however, the long-range interactions are not screened, then difficulties are encountered, such as the non-existence of the canonical ensemble\\cite{paddy} or inequivalence of microcanonical and canonical ensembles, and potential lack of a stable equilibrium configuration\\cite{antanov,lynden1}. The latter is sometimes attributed to negative heat-capacities\\cite{thirring}. Negative heat capacities are not only present in astrophysical systems, but have even been observed in fragmenting nuclei\\cite{negcap-nuc} and atomic clusters\\cite{negcap-sod}. It is not known how to deal with these systems generically, although there have been some attempts to understand them outside of standard thermodynamics using the Tsallis entropy\\cite{tsallis} [c.f. also \\cite{renyi}].     A principle motivation for this work is therefore to provide a framework in which to study such systems.  A second motivation comes from the study of black-hole thermodynamics.  There, it is found that the black-hole posseses an entropy which has unusual properties.  Here, we will show that these properties are not limited to the black-hole, but that other systems with long-range interactions exhibit related behavior.  We will essentially construct an analog of a black-hole by adding long-range interactions to a spin-lattice model.  Systems with long-range interactions are often referred to as {\\it non-extensive}, because the entropy and energy do not scale with the volume of the system. Normally, if one has a thermodynamical system, and one holds the intensive variables (temperature, pressure, and chemical potential) fixed, then if the size of the system is doubled, the extensive variables (entropy and energy) will also double.  This is not true if the interactions are long-range.  The purpose of this article is to develop new methods and a formalism to explore a number of facets of such non-extensive systems. We shall employ a principle of \"local-extensivity\" which enables one to define thermodynamical quantities for an interacting system.  We shall also show how to classify the degree of non-extensivity of the system, by calculating the scaling behavior of the entropy as a function of total energy. The motivation for this part of the study comes from general relativity. There, it is found that the entropy of a black-hole is proportional to its area, rather than its volume (i.e. the entropy is non-extensive).  In this study, we will see that this is a generic property of interacting systems, rather that something unique to the black-hole.  We will also see that generically, a system interacting with a reservoir is not at a particular temperature, but rather, is in a probability distribution of temperatures. This will be found by studying a system interacting with a reservoir in the microcanonical ensemble i.e. the total energy of the system plus reservoir is fixed. Usually, if one then only looks at the system, it will be in a canonical ensemble (fixed temperature). When interactions are present, this will not be the case. This leads us to introduce a new type of ensemble, which we call the microlocal ensemble.  It is equivalent to the microcanonical ensemble when there are no interactions.  Despite the fact that a system is not found in the canonical ensemble, we shall see that one can define a quantity we call the {\\it global temperature} $\\beta_o$.  It describes the system as a whole, and is written in terms of the total energy, including the interacting terms. There is also a {\\it local temperatures}, a quantity inspired by general relativity.  Both types of temperature are measurable in principle.   In the case of short-range interactions, two systems brought into thermal contact will be at the same temperature. Here, we will show that also for interacting systems, the global temperature of two systems is the same. This justifies to some extent, our use of the term temperature to describe $\\beta_o$.  However, we will see that if one has two separated systems with the same global temperature, then when they are brought into contact adiabatically, they will reach a new global temperature.  $\\beta_o$ is therefore not an intensive quantity.  We will also see that the local temperatures of the two systems are in general different for each system when in thermal equilibrium. Such an effect is analogous to the Tolman relation\\cite{tolman} which exists in curved space. There, one finds a temperature gradient due to the curvature of space-time. Essentially, frequencies are red-shifted by curvature.  Since the temperature gives the probability distribution of a frequency spectrum, it is also red-shifted. Here, we see that such an effect does not exist solely in curved space-time, but can also be thought of as due to the presence of long-range interactions.  To make our discussion more concrete, we will examine a toy model consisting of a lattice of spins in a magnetic field, and interacting via a spin-spin coupling.  However, rather than only nearest neighbor interactions, we will also consider the long-range couplings.  We will consider the case of a uniform long-range interaction, as well as the case of two different systems interacting via two unequal uniform interactions.  Such a situation arises when one considers two lattice clusters which are of small size. We will also discuss the continuum situation, where the interaction can be arbitrary, and varies from site to site.  We then consider thermodynamics in the general theory of relativity.  Comparisons between the lattice model and of black-hole thermodynamics provide another motivation for this study.  Previously, analogs of black-holes\\cite{dumbhole}\\cite{visser}, (so called, acoustic, or solid state black-holes) have been used to understand black-hole radiation. However, they are not useful for understanding the black-hole entropy. Here, we see that one can construct a a black-hole analog that can be used to study black-hole entropy. One finds that the entropy can be non-extensive, just like a real black-hole. We find for the analog that there is an infinite red-shifting between its local temperature and its global temperature which has exactly the same form as a black-hole. At exactly the point where the systems acts like a black-hole, a degeneracy in the local energy levels forms.  This degeneracy is universal, in the sense that it only depends on the form of the interaction. The universality is somewhat reminiscent of the universality of black-hole entropy.  We will also investigate other gravitational systems in general relativity. In particular, we look at the entropy scaling behavior of a gravitating perfect fluid.  The motivation for this comes partly from an earlier study\\cite{area} where I showed that the black-hole is not the only gravitating system which has an area-scaling entropy. A system of shells has an entropy that scales as the volume when the gravitational interaction is weak, but the entropy becomes area scaling at the point before a black-hole is formed.  Here, in looking at the gravitating perfect fluid, we find related behavior. We can look at the entropy scaling behavior not just in limiting cases, but for all strengths of the gravitational interaction. As the strength of the gravitational interaction is increased, the entropy slowly moves to the outer surface of the perfect fluid One finds that the total entropy is non-extensive, just like in a black-hole, and approaches area scaling behavior as the strength of the gravitational interaction gets stronger.  We also explore gravitating systems in the context of Newtonian-type dynamics. This is done to show that the red-shifting of temperatures -- usually considered to be an effect related to the curvature of geometry --  also exists in other gravitational models which are not geometric theories.  The paper is organized as follows.  In Section \\ref{sec:longrange} we introduce our formalism.  First, in Subsection \\ref{ss:localextensivity} we introduce our assumptions, which we call {\\it locality} and {\\it local extensivity} and show that these assumptions are obeyed by a number of common systems.  Next, in Subsection \\ref{ss:temp} we use the microcanonical ensemble to show that a system interacting with a reservoir will not be found at a particular temperature, but rather, will be in a probability distribution of different temperatures. Nonetheless, there is a parameter which behaves very similarly to a temperature, which we call the global temperature.  This is defined in Subsection \\ref{ss:global}.  The physical significance of the global temperature, as well as another parameter called the local temperature is explored in Subsections \\ref{ss:physsigofglobal} and \\ref{ss:physsigoflocal}.  Then, in Subsection \\ref{ss:unequal} we show that the global temperatures of two systems brought into contact, are equal at equilibrium.  The local temperatures need not be equal (an effect reminiscent of red-shifting which is usually considered to be the sole domain of general relativity).  This allows us to study lattice models where the long-range interaction is not uniform. Next, in Section \\ref{sec:exising} we explore in some detail a lattice model with long range interactions.  In Section \\ref{sec:bhanalog} we show that such a system can be made into an analog of a black-hole and has a temperature and entropy with  many properties reminiscent of black-holes. In Section \\ref{sec:scaling} we show how one can generically calculate the entropy scaling behavior of an interacting system.  This is done for a gravitating perfect fluid in Section \\ref{sec:entropyfluid}.  We find that the entropy becomes more area scaling as the gravitational interaction gets stronger.   We conclude with some general remarks in \\ref{sec:conclusion} and point to some open questions.  In Appendix \\ref{ap:newtonian} we look at systems under the influence of Newtonian-type gravity, and show that an analog of the Tolman relation exists -- local temperatures are red-shifted. In Appendix \\ref{ap:continuum} we look at more general interactions and go to the continuum limit. ",
        "conclusions": ""
    },
    "0212/astro-ph0212217_arXiv.txt": {
        "abstract": "s#1#2#3{{ \\centering{\\begin{minipage}{4.5in}\\footnotesize\\baselineskip=10pt \\parindent=0pt #1\\par \\parindent=15pt #2\\par \\parindent=15pt #3 \\end{minipage}}\\par}}  \\def\\keywords#1{{ \\centering{\\begin{minipage}{4.5in}\\footnotesize\\baselineskip=10pt {\\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}}  \\newenvironment{romanlist} {\\setcounter{romanlistc}{0} \\begin{list}{$($\\roman{romanlistc}$)$} {\\usecounter{romanlistc} \\setlength{\\parsep}{0pt} \\setlength{\\itemsep}{0pt}}}{\\end{list}}  \\newenvironment{alphlist} {\\setcounter{alphlistc}{0} \\begin{list}{$($\\alph{alphlistc}$)$} {\\usecounter{alphlistc} \\setlength{\\parsep}{0pt} \\setlength{\\itemsep}{0pt}}}{\\end{list}}  \\newenvironment{arabiclist} {\\setcounter{arabiclistc}{0} \\begin{list}{\\arabic{arabiclistc}} {\\usecounter{arabiclistc} \\setlength{\\parsep}{0pt} \\setlength{\\itemsep}{0pt}}}{\\end{list}}  \\newcommand{\\fcaption}[1]{ \\refstepcounter{figure} \\setbox\\@tempboxa = \\hbox{\\footnotesize Fig.~\\thefigure. #1} \\ifdim \\wd\\@tempboxa > 5in {\\begin{center} \\parbox{5in}{\\footnotesize\\smalllineskip Fig.~\\thefigure. #1} \\end{center}} \\else {\\begin{center} {\\footnotesize Fig.~\\thefigure. #1} \\end{center}} \\fi}  \\newcommand{\\tcaption}[1]{ \\refstepcounter{table} \\setbox\\@tempboxa = \\hbox{\\footnotesize Table~\\thetable. #1} \\ifdim \\wd\\@tempboxa > 5in {\\begin{center} \\parbox{5in}{\\footnotesize\\smalllineskip Table~\\thetable. #1} \\end{center}} \\else {\\begin{center} {\\footnotesize Table~\\thetable. #1} \\end{center}} \\fi}  \\def\\@citex[#1]#2{\\if@filesw\\immediate\\write\\@auxout {\\string\\citation{#2}}\\fi \\def\\@citea{}\\@cite{\\@for\\@citeb:=#2\\do {\\@citea\\def\\@citea{,}\\@ifundefined {b@\\@citeb}{{\\bf ?}\\@warning {Citation `\\@citeb' on page \\thepage \\space undefined}} {\\csname b@\\@citeb\\endcsname}}}{#1}}  \\newif\\if@cghi \\def\\cite{\\@cghitrue\\@ifnextchar [{\\@tempswatrue \\@citex}{\\@tempswafalse\\@citex[]}} \\def\\citelow{\\@cghifalse\\@ifnextchar [{\\@tempswatrue \\@citex}{\\@tempswafalse\\@citex[]}} \\def\\@cite#1#2{{$\\null^{#1}$\\if@tempswa\\typeout {IJCGA warning: optional citation argument ignored: `#2'} \\fi}} \\newcommand{\\citeup}{\\cite}  \\def\\pmb#1{\\setbox0=\\hbox{#1} \\kern-.025em\\copy0\\kern-\\wd0 \\kern.05em\\copy0\\kern-\\wd0 \\kern-.025em\\raise.0433em\\box0} \\def\\mbi#1{{\\pmb{\\mbox{\\scriptsize ${#1}$}}}} \\def\\mbr#1{{\\pmb{\\mbox{\\scriptsize{#1}}}}}  \\def\\fnm#1{$^{\\mbox{\\scriptsize #1}}$} \\def\\fnt#1#2{\\footnotetext{\\kern-.3em {$^{\\mbox{\\scriptsize #1}}$}{#2}}}  \\def\\fpage#1{\\begingroup \\voffset=.3in \\thispagestyle{empty}\\begin{table}[b]\\centerline{\\footnotesize #1} \\end{table}\\endgroup}   \\def\\thefootnote{\\fnsymbol{footnote}} \\def\\@makefnmark{\\hbox to 0pt{$^{\\@thefnmark}$\\hss}}\t%  \\def\\ps@myheadings{% \\let\\@oddfoot\\@empty\\let\\@evenfoot\\@empty \\def\\@evenhead{\\slshape\\leftmark\\hfil}% \\def\\@oddhead{\\hfil{\\slshape\\rightmark}}% \\let\\@mkboth\\@gobbletwo \\lets{ Using two different pseudo-Schwarzschild potentials proposed by Artemova et. al,$^{1}$ we formulate and solve the equations governing spherically symmetric transonic inflow and outflow in presence of a relativistic hadronic pressure mediated steady, standing, spherical shock around the central compact object and then we self-consistently connect the accretion-wind solutions to calculate the mass outflow rate $R_{\\dot m}$ in terms of minimum number of flow parameters. Also we study the dependence of this rate on various boundary conditions governing the flow. }{}{}   \\vspace*{1pt}\\textlineskip\t% \\vskip 0.5cm \\noindent \\hrule \\noindent {\\bf Published in International Journal of Modern Physics D, 2002, Volume 11, Issue 08, pp. 1285-1303.} \\hrule \\vskip 0.25cm ",
        "introduction": "% \\noindent For spherically symmetric accretion onto a non-rotating compact object, a novel mechanism was proposed by Protheros and Kazanas,$^{2}$ (PK83 hereafter) and Kazanas and Ellison$^{3}$ (KE86 hereafter), in which the kinetic energy of the accreting material was shown to be randomized by incorporating a steady, collision-less, relativistic hadronic-pressure-supported spherical shock surface around the accreting Schwarzschild black hole which produces a non-thermal spectrum of relativistic protons. Recently Das$^{4}$ (D99 hereafter) has explicitly calculated the exact location (radial distance measured from the central accretor in units of Schwarzschild radius $r_g=\\frac{2GM_{BH}}{c^2}$, where $M_{BH}$ is the mass of the black hole, $G$ is the Universal gravitational constant and $c$ is the velocity of light in vacuum) of the above mentioned shock in terms of only three accretion parameters, namely, the specific energy of the flow ${\\cal E}$, accretion rate ${\\dot M}_{Edd}$ (scaled in units of Eddington rate) and the adiabatic index $\\gamma_{in}$ of the inflow for spherically-symmetric, transonic accretion of adiabatic fluid onto a Schwarzschild black hole. By solving the set of hydrodynamic equations describing the motion of accreting material under the influence of modified Newtonian potential proposed by Paczy\\'nski and Wiita,$^5$ it has been shown there in D99 that it is possible to construct a self-consistent inflow-outflow system where a part of the accreting material may be blown as wind from the spherical shock surface and the mass outflow rate $R_{\\dot m}$ (the measure of the fraction of accreting material being `kicked off' as wind) was computed in terms of various accretion as well as shock parameters.\\\\ However, other than the Paczy\\'nski-Wiita potential, a number of modified Newtonian potentials of various forms are also available in the literature which accurately approximate some general relativistic features of rotating accretion around  Schwarzschild black holes. Such potentials may be called `pseudo - Schwarzschild' potentials because they mimic the space-time around a non-rotating/slowly rotating compact object. Recently Das and  Sarkar$^6$ (DS hereafter) has shown that though these so called pseudo potentials were originally proposed to mimic the relativistic effects manifested in disc accretion, it is quite reasonable to use most of these potentials in studying various dynamical and thermodynamic quantities also for spherical accretion on to Schwarzschild black holes and it was established that along with the Paczy\\'nski and Wiita potential,$^5$ two other potentials proposed by Artemova et. al.$^1$ also provide reasonably good approximation to the complete general relativistic description of transonic, spherically symmetric accretion on to a Schwarzschild black hole.\\\\ \\noindent Remembering  that the free-fall acceleration plays a very crucial role in Newtonian gravity, Artemova et. al.$^1$  proposed the following two potentials to study disc accretion around a non-rotating black hole. The first potential proposed by them produces exactly the same value of the free-fall acceleration of a test particle at a given value of $r$ as is obtained for a test particle at rest with respect to the Schwarzschild reference frame, and is given by $$ \\Phi_{1}=-1+{\\left(1-\\frac{1}{r}\\right)}^{\\frac{1}{2}} \\eqno{(1a)} $$ The second one gives the value of the free fall acceleration that is equal to the value of the covariant component of the three dimensional free-fall acceleration vector of a test particle that is at rest in the Schwarzschild reference frame and is given by $$ \\Phi_{2}=\\frac{1}{2}ln{\\left(1-\\frac{1}{r}\\right)} \\eqno{(1b)} $$ For Eq. 1(a-b), $r$ represents the radial co-ordinate scaled in the unit of $r_g$. Efficiencies produced by $\\Phi_1$ and $\\Phi_2$ are $-0.081$ and $-0.078$ respectively. As both $\\Phi_1$ and $\\Phi_2$ stems from the consideration of free fall acceleration and calculates the dependence of free fall acceleration on radial distance in the Schwarzschild metric (which describes a spherically symmetric gravitational field in vacuum), it appears to be quite justified to use those potentials to study spherically symmetric accretion.\\\\ \\noindent Owing to the fact that $\\Phi_1$ and $\\Phi_2$ may be used to mimic the spacetime around a spherically accreting non-rotating black hole quite nicely, we believe that along with the calculation presented in D99, it is equally important to investigate the various features of the accretion powered spherical outflow using these two potentials. In this paper we precisely do this, first we formulate and solve the equations governing the inflow and outflow using  $\\Phi_1$ and $\\Phi_2$ (in presence of a steady, standing spherical shock as described in previous paragraphs) and then self-consistently connect the accretion and wind solutions to calculate what fraction of the accreting material is being blown as wind. Also we study the dependence of this fraction on various flow parameters. The plan of this paper is as follows. In next section we describe how to formulate and solve the governing equations. In \\S 3, we present our results. In \\S 4, the results will be reviewed and conclusions will be drawn.\\\\ ",
        "conclusions": "\\noindent In this paper, we could successfully construct a self-consistent spherically-symmetric, polytropic, transonic, non-magnetized inflow-outflow system by simultaneously solving the set of hydrodynamic equations governing the accretion and wind around a Schwarzschild black hole using two different pseudo-potentials proposed by Artemova et. al.$^1$ Introducing a steady, standing, hadronic-pressure supported spherical shock surface around the black hole as the effective physical atmosphere which may be responsible for generation of accretion-powered spherical wind, we calculate the mass-outflow rate $R_{\\dot m}$ in terms of only three accretion parameters (conserved energy of the flow ${\\cal E}$, accretion rate ${\\dot M}_{Edd}$ scaled in units of Eddington rate and polytropic index of the flow $\\gamma_{in}$) and only one outflow parameter (the polytropic index of the outflow, $\\gamma_o$). Not only do we provide a sufficiently plausible estimation of $R_{\\dot m}$, we could also successfully study the dependence and variation of this rate on various physical parameters governing the flow.\\\\ \\noindent The basic conclusions of this paper may be summarized as: \\begin{enumerate} \\item Shock formation is not a generic phenomena, i.e., not all solutions contain shock, rather a specific region of parameter space spanned by ${\\cal E}, {\\dot M}_{Edd}$ and $\\gamma_{in}$ allows shock formation. For given values of ${\\cal E}, {\\dot M}_{Edd}$ and $\\gamma_{in}$, while the value of shock location (in units of $r_g$) {\\it correlates} with ${\\dot M}_{Edd}$, it {\\it anti-correlates} with both ${\\cal E}$ and $\\gamma_{in}$. High energy high luminosity purely non-relativistic accretion is a proper choice to produce high post-shock proton temperature( thus to produce the maximum outflow). For other accretion parameters being the same, $\\Phi_1$ produces a shock closer to the black hole compared to $\\Phi_2$. This deviation is more prominent for high ${\\cal E}$ and $\\gamma_{in}$. \\item The shock surface can serve as the `effective' physical barrier around the black hole regarding generation of mass loss via transonic spherical wind. The fraction of accreting material being blown as wind (which is denoted as $R_{\\dot m}$) could be computed in terms of three accretion parameters and one outflow parameter. \\item $R_{\\dot m}$ correlates with ${\\cal E}, {\\dot M}_{Edd}, \\gamma_{in}$ and $\\gamma_o$, Outflow could be generated for {\\it both} sub-Eddington as well as super-Eddington accretion. For a fixed set of \\\\ \\noindent $\\left\\{{\\cal E}, {\\dot M}_{Edd}, \\gamma_{in}, \\gamma_o\\right\\}$, mass outflow rate $R_{\\dot m}$ is higher for flows in $\\Phi_2$ compared to flows in $\\Phi_1$, which indicates that the barionic content of the spherical wind is inversely proportional to the spatial gradient on the pseudo-potential used to study the problem though the exact physical reason behind this is not quite clear. \\item If a shock forms, then whatever the initial flow conditions and whatever the nature of dependence of $R_{\\dot m}$ on any of the accretion/ shock/ outflow parameters, $R_{\\dot m}$ normally correlates with post-shock flow temperature, which indicates that outflow is strongly thermally driven; hotter flow always produces more winds.\\\\ \\end{enumerate} \\noindent At this point, it is worth mentioning that the hot and dense shock surface around black holes, which is proposed here as the effective physical barrier around compact objects regarding the mass outflow, may be generated due to other physical effects as well for spherical accretion.$^{10,11,12,13,14}$ Another very important approach launched recently was to construct such an `effective barrier' for non-spherical disc accretion to introduce the concept of CENtrifugal pressure supported BOundary Layers (CENBOL). Treating the CENBOL as the effective atmosphere of the rotating flows around compact objects (which forms as a result of standing Rankine-Hugoniot shock or due to the maximization of polytropic pressure of accreting material in absence of shock), detailed computation of the mass-outflow rate from the advective accretion disks has been done, and dependence of this rate on various accretion and shock parameters has been quantitatively studied by constructing a self-consistent disk-outflow system.$^{15,16}$ \\\\ Our calculations in this paper, being simply founded, do not explicitly include various radiation losses and cooling processes, combined effects of which may reduce the post-shock proton temperature (which means the reduction of outflow temperature), in reality could be lower than what we have obtained here and the amount of outflow would be less than what is obtained in our calculation. This deviation will be more important for systems with high accretion rates. Nevertheless, cases of low accretion rates discussed here would not be affected that much and our preliminary investigation shows that even if we incorporate various losses, the overall profile of the various curves showing the dependence of $R_{\\dot m}$ on different inflow parameters would be exactly the same, only the numerical value of $R_{\\dot m}$ in some cases (especially for high accretion) might decrease.\\\\ \\noindent Although in this work we have performed our calculation for a 10$M_{\\odot}$ Schwarzschild black hole, general flow characteristics will be unchanged for black hole of any mass except the fact that the region of parameter space responsible for shock formation will be shifted and the value of $R_{\\dot m}$ will explicitly depend on the mass of the black hole.\\\\[0.25cm] \\nonumsection{Acknowledgements} \\vspace*{-0.5pt} \\noindent The author would like to thank Prof. I. Novikov for useful discussion which enabled him to have a better understanding about the physical nature of the pseudo-potentials used in this work.\\\\ [0.25cm] \\vspace*{-0.5pt} \\nonumsection{References} \\vspace*{-0.5pt}"
    },
    "0212/astro-ph0212203_arXiv.txt": {
        "abstract": "{The question is answered whether $\\alpha^2$-shell-dynamos are able to produce a cyclic activity or not. Only  kinematic dynamos are considered and  only the solutions with the lowest dynamo number are studied without restrictions  about the axial symmetry of the solution. The \\alf-effect is allowed to be latitudinally inhomogeneous and/or anisotropic, but it is assumed as radially uniform in the turbulent shell.\\\\ For a symmetric  \\alf-tensor we only find oscillatory solutions if three conditions are simultaneously fulfilled: i) the $\\alpha_{zz}$ vanishes or is of the opposite sign as $\\alpha_{\\phi\\phi}$, ii) the \\alf-effect is strongly concentrated to the equatorial region (i.e. it vanishes at the poles) and iii) the \\alf-effect is concentrated to a  rather thin outer shell. In the  other cases almost always the nonaxisymmetric field mode S1 possesses  the lowest dynamo number which slowly drifts along the azimuthal direction.  Also uniform but anisotropic \\alf-effect ($\\alpha_{zz} = 0$) leads to the nonaxisymmetric solutions  as it is  confirmed by the Karlsruhe dynamo experiment. \\\\ However, one of the antisymmetric parts of the \\alf-tensor ({\\em not} the vertical magnetic pumping) basically  plays the role of a differential rotation in the induction equation. Using for the radial profile of this effect the results of a numerical simulation for the \\alf-tensor of the solar convection zone (Ossendrijver et al. 2002), one indeed finds the possibility of oscillating $\\alpha^2$-dynamos even without the existence of real  nonuniform plasma rotation so that they really can be called as  pseudo-\\alf$\\Om$-dynamos. The resulting butterfly diagram, however, proves to be  of the antisolar type. The radial gradient of this pseudo-differential rotation determines the sign of the phase relation of $B_r$ and $B_\\phi$ (see Sect. \\ref{sec4}). ",
        "introduction": "There are two observations which in the past led to a increasing interest in the solutions of the relatively simple $\\alpha^2$-dynamo. The first one is the cyclic orbital modulation of close binary systems such as reported by Hall (1990) and Lanza \\& Rodon\\`{o} (1999) for RSCVn stars and the second one is the flip-flop phenomenon as reported recently by Tuominen et al. (1999) and Korhonen et al. (2001) for FK Coma stars but also in the single young dwarf LQ Hya (K2V, $P_{\\rm rot} = 1.6$ days, see Rice \\& Strassmeier 1998). Together with the highly nonaxisymmetric field configurations for very young cool dwarf  stars reported by Jardine et al. (2002) (see Fig. \\ref{fff1}) there is increasing evidence that the traditional stationary and axisymmetric dipole-solution which is known since decades does not form the final truth. \\begin{figure} \\hbox{\\hspace{1cm} \\psfig{figure=95_180_pole.ps,width=6cm,height=8cm} } \\caption{Pole-on representation of the  magnetic field of AB Dor. Courtesy M. Jardine.} \\label{fff1} \\end{figure}  Oscillating $\\alpha^2$-dynamos have been found to occur in the special circumstance wherein the $\\alpha$-effect changes rapidly in boundary layers (R\\\"adler \\& Br\\\"auer 1987; Baryshnikova \\& Shukurov 1987). In such cases, the period of the $\\alpha^2$-dynamo depends strongly on the location of the $\\alpha$-boundary layer and is typically an order of magnitude or more smaller than the magnetic diffusion time across the dynamo-generation shell. In general, oscillatory dynamo behavior is produced by a combination of both the $\\alpha$-effect and the $\\Om$-effect. Kinematic models of the oscillatory solar dynamo are always of the $\\alpha\\Om$-type. Here, also motivated by the publication of Schubert \\& Zhang (2000), we discuss the question whether also $\\alpha^2$-dynamos alone are able to exhibit oscillatory solutions. Schubert \\& Zhang (2000) considered $\\alpha^2$-dynamos with the (academic) assumptions of completely homogeneous and isotropic \\alf-effect and found oscillatory solutions for thin outer shells. In Section \\ref{3.2} we shall rediscuss this constellation under the extra condition that only the lowest eigenvalue provides a stable solution. Then all solutions prove to be non-oscillating.  The present paper only concerns kinematic dynamos on the basis of \\alf-tensors derived by numerical simulations of MHD-convection. Nonlinearities in the mean-field dynamo equations are not considered. They are highly complicated as the majority of the resulting magnetic field configurations is nonaxisymmetric (see Moss \\& Brandenburg 1995; K\\\"uker \\& R\\\"udiger 1999).  Within the considered turbulent shell the $\\alpha$-effect is assumed as uniform in radius, there is no change of its sign. Stefani \\& Gerbeth (2002) have shown that magnetic oscillations can occur if two shells with different signs of the $\\alpha$-effect exist.\\footnote{The oscillating solution, however, seems to disappear for geometries where one sign of the $\\alpha$-effect dominates the other as it is the case in Fig. \\ref{ossen1}.} We do not follow this possibility in the present paper, here the $\\alpha$-effect has been considered  as uniform in radius within the turbulent shell. ",
        "conclusions": "We have shown that for spherical configurations with outer turbulent shells the basic solution for $\\alpha^2$-dynamos is a nonaxisymmetric mode drifting in the azimuthal direction. Oscillating axisymmetric solutions are seldom exceptions for \\alf-effects that are \\begin{itemize} \\item[i)] highly anisotropic, \\item[ii)] strongly concentrated to the equatorial region, \\item[iii)] restricted to thin outer shells. \\end{itemize} In all other cases the mode with the lowest eigenvalue (which are considered here as the only stable one) is nonaxisymmetric and almost always of quadrupolar symmetry. The azimuthal drift is always of the same order as the oscillation frequencies are. Note that the box simulations by Ossendrijver et al. (2001) did {\\em not} lead to vanishing \\alf-effect at the poles.  The same is true for $\\alpha^2$-dynamos under inclusion of those anisotropic parts of the \\alf-tensor which can be combined to antisymmetric components, i.e. $G_i {\\Om}_m - G_m {\\Om}_i$. This term corresponds to a (pseudo-) angular velocity, see Eq. (\\ref{omt}), which induces electrical fields in the same way as a real differential rotation of the plasma would do. The box simulations of Ossendrijver et al. reveal this pseudo-${\\Om}$ (called here ${\\Om}_{\\rm T}$) as increasing outwards (see Fig. \\ref{ossen2}) so that in any case, together with the positive $\\alpha_{\\phi\\phi}$ (in the northern hemisphere), an antisolar butterfly diagram results if the solution was axisymmetric and oscillating. The necessary dipolar solution, however,  only exists if the ${\\Om}_{\\rm T}$-effect is artifically amplified by a factor of 10 (see Table \\ref{tab01}).  We conclude that shell-dynamos without differential rotation on the basis of MHD simulations  of rotating stellar convection always possess nonaxisymmetric magnetic field configurations such as recently found for AB Dor.  Oscillating solutions for \\alf$^2$-dynamos are revealed as rather exceptions. Our conclusion is that a cyclic stellar activity can always be considered as a strong indication for the existence of internal differential rotation."
    },
    "0212/astro-ph0212035_arXiv.txt": {
        "abstract": "{ The 13hr {\\em XMM-Newton}/{\\em Chandra} deep survey is the first of two extremely deep {\\em XMM-Newton} fields observed by the XMM-OM consortium. A 120 ks {\\em Chandra} mosaic, covering 0.2~deg$^{2}$, provides sensitive, confusion-free point source detection with sub-arcsecond positions, while the 200 ks {\\em XMM-Newton} observation provides high quality X-ray spectroscopy over the same sky area. We have optical spectroscopic identifications for 70 X--ray sources. Of these, 42 are broad emission-line AGN with a wide range of redshifts. The optical counterparts of a further 23 sources are narrow emission line galaxies and absorption line galaxies. These 23 sources all lie at $z<1$ and typically have lower X--ray luminosities than the broad-line AGN. About half of them show significant X--ray absorption and are almost certainly intrinsically absorbed AGN. However some of them have unabsorbed, AGN-like, power-law components in their X--ray spectra, but do not show broad emission lines in their optical spectra. These sources may be weak, unobscured AGN in bright galaxies and their existence at low redshifts could be a consequence of the strong cosmological evolution of AGN characteristic luminosities. ",
        "introduction": "Above 1 keV, the cosmic X--ray background (CXRB) is predominantly the integrated emission of many discrete sources. In the 0.5-2 keV X--ray band, the deepest surveys with {\\em ROSAT} resolved 70-80\\% of the X--ray background into individual sources, the majority of which are broad-line AGN (e.g.  Hasinger et al. \\nocite{hasinger98} 1998, McHardy et al. \\nocite{mchardy98} 1998). However the energy density of the CXRB peaks at much higher energy, $\\sim 30$ keV, and to reproduce its broadband spectral shape, CXRB synthesis models require a large population of hard-spectrum X--ray sources. The most enduring of these synthesis models invoke a large population of absorbed AGN, as expected in AGN `unification schemes' (Setti and Woltjer \\nocite{setti89} 1989, Comastri et al. \\nocite{comastri95} 1995, Gilli, Risaliti \\& Salvati \\nocite{gilli99} 1999). Deep surveys performed with {\\em XMM-Newton} and {\\em Chandra} have now succeeded in resolving a large fraction of the CXRB up to 10 keV (e.g. Brandt et al., \\nocite{brandt01} 2001, Campana et al. \\nocite{campana01} 2001, Hasinger et al., \\nocite{hasinger01} 2001). The X-ray and optical properties of the {\\em Chandra} and {\\em XMM-Newton} sources will allow us to test these synthesis models.  The 13hr deep field is a project to investigate the astrophysics of the major contributors to the CXRB, particularly sources around the break in the X--ray source counts where the contribution to the CXRB per log-flux interval peaks. In order to understand the phenomena, processes and conditions in these sources we combine the high quality X--ray spectra of {\\em XMM-Newton} with the precise positions of {\\em Chandra} and a host of other multi-wavelength data. The 13hr field was the location of one of the deepest {\\em ROSAT} surveys (McHardy et al. \\nocite{mchardy98} 1998) and is a region of extremely low Galactic absorption (${\\rm N_{H}} \\sim 6.5 \\times 10^{19}$~cm$^{-2}$). ",
        "conclusions": "As seen in Fig. \\ref{fig:specresults}, the underlying spectral indices of the sources are scattered around the mean value of $\\Gamma \\sim 2$ with a standard deviation of $\\sim 0.5$, and show no obvious dependence on luminosity, absorption, or optical classification. Furthermore, the distribution of spectral indices is remarkably similar to that found in brighter, soft X--ray selected samples (e.g. Mittaz et al. \\nocite{mittaz99} 1999). Thus the {\\em XMM-Newton} spectroscopy is consistent with a picture in which the underlying power law X-ray spectra of AGN do not depend on redshift or luminosity, and are the same in the absorbed and unabsorbed sources.  In the context of the `standard' picture of the X-ray populations, the majority of the NELGs and galaxies, and certainly those with $L_{X} > 10^{42}$~erg~s$^{-1}$, are expected to be intrinsically absorbed AGN. As seen in Table \\ref{tab:IDs} and Fig. \\ref{fig:specresults}, absorption is detected in $\\sim 50\\%$ of the galaxies and NELGs compared to only $\\sim 10\\%$ of the BLAGN; these latter sources have $z > 1$ and column densities between $5\\times 10^{21}$~cm$^{-2}$ and $2\\times 10^{22}$~cm$^{-2}$ and so are similar to the luminous, X-ray absorbed, BLAGN found in the hard-spectrum {\\em ROSAT} survey of Page, Mittaz \\& Carrera \\nocite{page01} (2001). Thus a significant fraction of the NELGs and galaxies are indeed absorbed AGN, as reported for the Lockman Hole (Mainieri et al. \\nocite{mainieri02} 2002).  However, a significant fraction ($\\sim 50\\%$) of the galaxies and NELGs do not show strong evidence for X--ray absorption, but have power-law (i.e. AGN-like) X--ray spectra and the majority of them have AGN X--ray luminosities.  The lack of broad emission lines in these sources is not readily explained as an absorption effect: assuming a Galactic dust to gas ratio, a column density of $5 \\times 10^{21}$ cm$^{-2}$ is required to attenuate broad lines such as H$\\beta$ by a factor of 10. Such column densities are readily detectable for the (mostly low-$z$) galaxies and NELGs.  \\begin{figure} \\resizebox{\\hsize}{!} {\\includegraphics[]{chandra59.eps}} \\caption{Keck LRIS red and blue arm spectra of Chandra~59, a BLAGN with z=0.546. In the red, the spectrum is dominated by the bright host galaxy but in the blue (rest-frame UV) the AGN dominates and the broad Mg~II emission line is obvious} \\label{fig:chandra59} \\end{figure}  One feature that is noticeable in Fig. \\ref{fig:specresults} is that the X--ray sources classified as NELGs and galaxies typically have lower X--ray luminosities than the BLAGN. This is true for both the X--ray absorbed and unabsorbed NELGs and galaxies, even after corrections for absorption: {\\it none} of the NELGs or galaxies have log~$L_{X} > 43.7$, but the {\\it majority} of the BLAGN have log~$L_{X} > 43.7$. A likely explanation for the lack of broad emission lines in the unabsorbed galaxy and NELG sources is therefore simply one of contrast: these sources are weak (e.g. factor 10 underluminous) AGN in bright host galaxies. In fact even amongst the sources which {\\em are} classified as BLAGN, we can see this kind of phenomenon: Fig. \\ref{fig:chandra59} shows one such source, which is only recognisable as a BLAGN in the rest frame UV where the host galaxy is weak. Thus the optical classification of the lower luminosity AGN is likely to have a complex dependancy on both redshift and the host-galaxy luminosity. However, the presence of significant power law components in the X--ray spectra of NELGs and galaxies demonstrates that they do contain AGN; this has important implications for the evolution of the AGN X--ray luminosity function (e.g. see Page et al. \\nocite{page97} 1997). Similar hypotheses of AGN/host-galaxy contrast have also been proposed to hide emission lines in Seyfert 2 galaxies (Lumsden \\& Alexander \\nocite{lumsden01} 2001; Moran, Filippenko \\& Chornock \\nocite{moran02} 2002). Another mechanism which can make broad emission lines undectable is beamed continuum emission from BL Lacs; there are at least two examples of this type of object in the 13hr field (see Gunn et al. this proceedings).  If the lack of broad lines in X--ray unabsorbed AGN is a consequence of the relative brightness of AGN and their host galaxies, the cosmological evolution of AGN may offer an explanation as to why a large fraction of NELGs and galaxies (and the equivalent sources in Rosati et al. \\nocite{rosati02} 2002 and Mainieri et al. \\nocite{mainieri02} 2002) are found predominantly at $z < 1$. The characteristic luminosity (i.e. the break in the luminosity function) of AGN has declined dramatically (by a factor $>10$) since $z=2$ at both X--ray (e.g. Page et al. \\nocite{page97} 1997) and optical wavelengths (e.g. Boyle et al. \\nocite{boyle00} 2000). Unless the host galaxies of AGN have declined in luminosity by a similar amount, the contrast between AGN and host galaxy will typically be smaller today than it was at high redshift."
    },
    "0212/astro-ph0212345_arXiv.txt": {
        "abstract": "All quasars vary  in their optical flux on a  time-scale of years, and it  has been  proposed that  these variations  are principally  due to gravitational lensing by  a cosmologically distributed population of planetary  mass objects.  This interpretation has  implications for the observable  properties of gamma-ray  bursts (GRBs) -- as  a source expands  across  the nano-arcsecond  caustic  network, variability  is expected -- and data on GRBs can be used to test the proposed model of quasar variability. Here we employ an ultra-relativistic blast-wave model of the source, with no intrinsic variations, to study the effects of nanolensing on GRBs. Taken in isolation the light-curves of the caustic crossings are predictable, and we find that a subset of the predicted light-curves (the image-annihilating  fold crossings) resemble the ``pulses'' which are commonly seen in long GRBs. Furthermore, for sources at high redshift the expected time between  caustic crossings is  of order  seconds,  comparable to  the observed time  between pulses. These  points suggest that it might be possible to model some of the observed variations of GRBs in terms of  nanolensing; however, our simulated light-curves  exhibit a small depth of  modulation compared to what is observed. This means that the GRB data do not significantly constrain the quasar nanolensing model; it also means that the simplest nanolensing model cannot explain the observed GRB ``pulses''. Viable nanolensing models for pulses probably require a large external beam shear.  If  a  viable model  can  be constructed it would effect  a considerable simplification in source modelling and,  ironically, it would explain why  no macro-lensed GRBs have been identified to date.  Independent of the particular theoretical model, we can test for the presence of nanolensing in GRB data because any  variability  due to  nanolensing  should  manifest parallax:  the timing of  caustic crossings, and  hence the temporal  substructure of bursts should  be different as seen by  separated observers.  Parallax therefore  shifts triangulated  burst locations  away from  their true positions; this  displacement is typically  expected to be at  the few arcminute level, and existing astrometry  is not good enough to reveal the predicted  effects.  Useful constraints can,  however, be obtained by  comparing   the  relative  timing  of  individual   peaks  in  the light-curves  recorded by spacecraft  in the  Inter-Planetary Network; published data  show hints of  the predicted temporal shifts,  but the photon counting statistics are not good enough to categorically decide the matter. There is  no plausible alternative interpretation for this phenomenon, and  if it is confirmed  as a real effect  then it compels acceptance of  a cosmology that  is very different from  the currently popular model. ",
        "introduction": "At present  we do not know  what constitutes the bulk  of the material Universe, i.e. the dark matter.  Many dark matter candidates have been proposed,  ranging from  elementary particles  to  macroscopic objects such as black holes --  see, e.g., Trimble (1987), Ashman (1992), Carr (1994). To  date it  has been possible  to eliminate  some suggestions (e.g.   brown dwarfs:  Tinney 1999),  by  virtue of  a clear  conflict between models  and data, but  a positive identification has  not been achieved. One proposal whose implications have not yet been thoroughly explored is that of Hawkins (1993,  1996), who argued, on the basis of photometric  monitoring  of  quasars,  that the  Universe  contains  a near-critical density  of planetary-mass  objects (cf. Press  and Gunn 1973).   Gravitational lensing  by  such a  population  is capable  of explaining much  of the observed  variability (see also  Schneider and Wei\\ss\\ 1987;  Schneider 1993) --- though  this should not  be taken to imply  that intrinsic  variations are  absent.  This  suggestion  is a radical departure from the now-standard picture of  a  universe  dominated  by  elementary  particles (i.e. Cold  Dark Matter -- see,  e.g., Peebles 1993;  Blumenthal et al 1984;  Davis  et  al   1985),  and  has  received  little  theoretical attention.   Building on  earlier  work (Schneider  and Wagoner  1987; Rauch 1991; Seljak and Holz 1999;  Metcalf and Silk 1999), Minty, Heavens  and Hawkins (2001) described  a  statistical  test  of  the  model  based  on  the light-curves  of distant supernovae,  but existing  data do  not allow these  ideas to  be  usefully  implemented. A  test  based on  surface brightness variability in low-redshift  galaxies has been described by Lewis and Ibata (2001), but this idea also requires data which are not yet available, as does the test based on monitoring quasars seen through low-redshift clusters of galaxies (Walker and Ireland 1995; Tadros, Warren and Hewett 1998). Data from the MACHO and EROS experiments exclude planetary-mass compact objects as a significant contributor to the Galactic dark matter (Alcock et al 1998). However objects which are sufficiently compact that they qualify as strong gravitational lenses at cosmological distances are not necessarily strong gravitational lenses when they are located in the Galactic halo (Walker 1999; see also Draine 1998, Rafikov and Draine 2001). The Galactic microlensing experiments therefore do not directly test the quasar nanolensing hypothesis.  It  is  obvious  that  one  could  attempt  to  devise  further,  more sophisticated tests  of the nanolensing interpretation based  on existing quasar data, but it is a priori unlikely that any such test will yield a definitive result. The main reason  for this is source size: the low amplitude of variations seen in quasar optical light-curves means that these  sources  are necessarily  larger  than  the  scale-size of  the hypothesised  caustic  structure,  thus  smoothing  the  magnification pattern,  and in  the  process destroying  the high  spatial-frequency information. (This would be less of  a barrier to progress if we had a reliable model for the physical structure of quasars, so that detailed predictions  for their  lensed  appearance could  be  given with  some confidence.) For  example the approach  of Dalcanton et al  (1994; see also Canizares 1982), based on  the statistics of equivalent widths of quasar emission lines,  is quite insensitive in this  regime where the continuum  source is  resolved  by  the lenses.  To  make progress  we therefore  require  a numerous  population  of small,  highly-luminous sources  at  large distances;  gamma-ray  bursters  constitute such  a population. Gravitational  lensing of gamma-ray  bursts has previously been  considered  by a  number  of  authors  (e.g. Paczy\\'nski  1986b; McBreen and  Metcalfe 1988;  Mao 1992; Gould  1992; Blaes  and Webster 1992; Nemiroff and Gould  1995), but their considerations are relevant to lenses  which are  either much more  massive, or much  less massive than the planetary-mass range which we consider here. The present work also differs  substantively from  previous investigations in  that the evolution (expansion) of the source structure plays a critical role in our analysis.  Following  the  launch of  the  Compton  Gamma-Ray Observatory,  BATSE (Burst And Transient Source Experiment) soon discovered that gamma-ray bursts are isotropically distributed on  the sky, and that they depart from Euclidean source counts at  low flux levels (Fishman et al 1994). These   discoveries   immediately    shifted   attention   away   from interpretations based on Galactic  neutron stars, which had previously been popular, towards a variety  of models in which energies amounting to  a significant  fraction of  $\\msun  c^2$ are  rapidly released  by sources at  cosmological distances (see  Paczy\\'nski 1995).  Discovery of X-ray afterglows from these  transient events (Costa et al 1997) by BeppoSAX allowed the sources to be accurately positioned, and thus led to the discovery  of optical and radio afterglows  (van~Paradijs et al 1997; Djorgovski et al 1997;  Frail et al 1997). Spectroscopy of these optical transients  has in some  cases revealed absorption  lines from gas  at redshifts  $z_s\\ga1$,  thus firmly  establishing the  distance scale  of  the  bursters  to  be cosmological  (Metzger  et  al  1997; van~Paradijs, Kouveliotou and Wijers 2000).  Although we still  do not know what process injects the burst energy, it is broadly agreed that the   radiation  observed from GRBs arises  from  a  relativistically  expanding source.  Relativistic  expansion is required  in order that  we should see  gamma-rays at  all  --  at least  at  energies $\\ga$~MeV  -- otherwise the inferred source-size is  so small, and the photon energy density so high,  that two-photon pair-production converts essentially all  the  gammas  into  material  particles (Cavallo  and  Rees  1978; Fenimore,  Epstein  and  Ho  1993;  Baring and  Harding  1997).  These considerations  require expansion  with  Lorentz factors  of at  least $\\Gamma\\sim10^2$.  At  an  early  stage  in  the  modelling  of  gamma-ray  bursts  in  a cosmological  context,  it was  pointed  out  that gamma-ray  emission should arise as the the  ambient medium undergoes shock compression by the expanding  material (Rees \\&  M\\'esz\\'aros 1992). This  picture -- the ``blast-wave'',  or ``external shock''  model -- now  provides the accepted  context for  modelling  the broad-band  (X-ray, optical  and radio) afterglows of bursts (Paczy\\'nski and Rhoads 1993; M\\'esz\\'aros and  Rees 1997;  Granot, Piran  and Sari  1999), but  has  fallen into disfavour  as an  interpretation of  the prompt  gamma-ray  burst. The reason  for the  demise of  this model  is simply  that it  cannot, in itself,  accommodate the rapid,  large-amplitude variability  which is manifest in  the gamma-ray data (Fenimore, Madras  and Nayashkin 1996; Sari  and Piran  1997;  see  also Dermer  and  Mitman 1999;  Fenimore, Ramirez-Ruiz and Wu 1999).  This  point has promoted the idea that the gamma-rays arise from internal shocks within the relativistic outflow, so that the  temporal variations of the burst  reflect the input power variations of the source  (e.g. Rees and M\\'esz\\'aros 1994; Kobayashi, Piran and Sari 1997).  Because  we   are  investigating   an hypothesis  in   which  apparent variability   arises  external  to   the  source,   as  a   result  of gravitational  lensing,  arguments  which  assume  that  the  observed variations are  intrinsic immediately lose  most, if not all  of their force. Given  this there are at  least three motivations  to return to the blast-wave interpretation of the prompt gamma-ray emission: first, this picture allows the initial  deposition of energy to be impulsive, and  is  therefore  suitable  for  a  broad  range  of  source  models (independent  of the  specific physics  of the  energy input)  with no fine-tuning  required;  secondly,  the  low  radiative  efficiency  of internal shocks  (Spada, Panaitescu and M\\'esz\\'aros  2000) appears to be inconsistent  with the event  energetics in at least  some cases (Paczy\\'nski 2001); and thirdly,  gamma-ray emission from  an external shock  ought  to be  present  at some  level.   This  emission can  be sensibly described  by a self-similar,  decelerating blast-wave model, similar to  those developed for  the afterglow emission  (e.g. Granot, Piran  and Sari  1999), but  radiative rather  than adiabatic.  Such a description leads  us to  expect a circular  source with  {\\it very\\/} strong limb brightening and consequently we   anticipate   significant flux variations if  the limb crosses a caustic.  This thin, bright and rapidly expanding ring is  a near-perfect instrument for revealing any caustic structure that might be present along the line-of-sight to the source --- see also Loeb and Perna (1998), and Mao and Loeb (2001), who considered microlensing of optical afterglows by stellar-mass objects in the low optical-depth limit.  In this  paper we use the  source model just described  to explore the hypothesis that the Universe contains a high density of planetary-mass objects,  with an  optical  depth to  gravitational  lensing of  order unity.  Often  the characteristic  angular  scale  of  the lenses  (in arcseconds)  is  indicated in  the  nomenclature  given to  associated phenomena (e.g.  {\\it micro\\/}lensing), and  following this convention dictates the  name ``nanolensing''  for the effects  discussed herein. Although  gravitational lensing  is the  main focus  of  our attention here, in  studying the  variations which might  be introduced  by this process we are  also implicitly addressing the physics  of the sources themselves. We start by presenting our source model in \\S2, and in \\S3 we turn  to aspects of  the gravitational lensing;  light-curves which result from the marriage of  these elements are presented and compared with data in \\S4. The role of parallax is explored in \\S5, followed by a  discussion of  related issues  concerning both  the lenses  and the sources (\\S6), and our conclusions are given in \\S7. ",
        "conclusions": "Motivated by an existing interpretation of quasar variability in terms of  gravitational nanolensing,  we have  examined the  implications of this model for  the observed properties of GRBs.  Using a self-similar blast-wave model to represent the source (no intrinsic variability) we find that the  light-curves of some of the  caustic crossings resemble those of the  ``pulses'' commonly seen in GRBs,  and for high redshift sources  the time-scale  of  the predicted  nanolensing variations  is consistent  with  the  GRB  data.  However,  the  predicted  depth  of nanolensing modulation is far too small to explain the deep variations observed in GRBs,  and this problem is exacerbated if  the GRBs are at high redshift. These results mean that the GRB data do not exclude the nanolensing  interpretation  of  quasar variability;  conversely,  the simplest (shear  free) nanolensing  model cannot explain  the observed GRB variability. Despite this failure, there  are weak  indications in the  published  InterPlanetary   Network  data  that  nanolensing  may actually be responsible for some of the observed  variations  of GRBs: the light-curves  for one GRB show  hints of parallax. This effect  is uniquely associated with lens-induced variations, and thus motivates a careful examination of existing IPN data."
    },
    "0212/astro-ph0212490_arXiv.txt": {
        "abstract": "{Optical emission-line profiles were evaluated in order to explore the structure of galactic nuclei containing H$_2$O megamaser sources. Long-slit spectra of IC\\,2560, NGC\\,1386, NGC\\,1052 and  Mrk\\,1210 were obtained at $\\sim 100$\\,km/s spectral and $\\sim 2\\arcsec \\times 2 \\arcsec$ spatial resolution. The following individual properties of the objects were found: The active nucleus of IC\\,2560 (innermost $\\pm 2\\arcsec$) emits lines typical for a high-ionization Seyfert-2 spectrum albeit with comparatively narrow profiles (FWHM $\\sim 200$\\,km/s). Line wings are stronger on the blue side than on the red side suggesting outflow. The central velocity gradient fits into the general velocity curve of the galaxy. Attributing it to a rotating disk coplanar with the galaxy leads to a Keplerian mass of $\\sim 10^7$\\,M$_{\\odot}$ inside a radius of 100 pc. --- The central 6\\arcsec~sized structure seen on HST H$\\alpha$ and [O{\\sc iii}] images of NGC\\,1386 appears to be the inner part of a near-edge-on warped rotating spiral disk that is traced in H$\\alpha$ within a diameter of 17\\arcsec~along p.a.\\,23\\degr. This interpretation is based on observed kinematic continuity and a typical S-shaped dust lane crossing the kinematical center. The central velocity gradient yields a Keplerian mass estimate of $\\sim$10$^8$\\,M$_{\\odot}$ inside $R=100$ pc and 5\\,10$^9$\\,M$_{\\odot}$ inside 0.8 kpc. The total mass of the ionized gas of $10^{5-6}$\\,M$_{\\odot}$ is small compared to the dynamical mass of the spiral disk. --- The kinematical gradient of the rotating gas disk in the center of the elliptical galaxy NGC\\,1052 yields a mass of $\\sim$6\\,10$^8$\\,M$_{\\odot}$ inside 166\\,pc. Two components are found: component $\\cal A$ arises in a rotating disk, component $\\cal B$ is blueshifted by $\\sim400$\\,km/s relative to the disk. $\\cal B$ likely originates from outflowing gas distributed within a wide cone. There is no need for a broad H$\\alpha$ component in unpolarized flux. The peculiar polarization seen in [O{\\sc i}]$\\lambda 6300$ by Barth et al. (1999) may be related to the outflow component. --- In Mrk\\,1210, a redshifted component $\\cal R$ contributes to the exceptional width of [O{\\sc i}]$\\lambda 6300$. An additional blueshifted component $\\cal B$ strong in [O{\\sc iii}] allows to fit the H$\\alpha$ + [N{\\sc ii}] blend without a broad component of H$\\alpha$. $\\cal B$ and $\\cal R$ are likely to be outflow components. The location of the brightness maximum BM is slightly ($\\sim1\\arcsec$) shifted relative to the kinematical center KC. Locating the true obscured nucleus in KC makes it understandable that BM displays faint broad scattered H$\\alpha$ in polarized light. --- Galactic rotation and outflow of narrow-line gas are common features of this sample of water-megamaser galaxies. All decomposed line-systems exhibit AGN typical line ratios. Recent detections of H$_2$O megamasers in starburst galaxies and the apparent asssociation of one megamaser with a Seyfert 1 AGN suggest that megamasers can possibly be triggered by optically detectable outflows. The frequently encountered edge-on geometry favoring large molecular column densities appears to be verified for NGC\\,1386 and IC\\,2560. For NGC\\,1052 and Mrk\\,1210, maser emission triggered by the optically detected outflow components cannot be ruled out. ",
        "introduction": "During the past two decades, powerful H$_2$O masers have been detected in the 6$_{16}$--5$_{23}$ transition at 22\\,GHz ($\\lambda$$\\sim$1.3\\,cm) towards almost thirty galaxies known to contain an active galactic nucleus (AGN) of Seyfert 2 or LINER type (dos Santos \\& L{\\'e}pine 1979; Gardner \\& Whiteoak 1982; Claussen et al. 1984; Henkel et al. 1984, 2002; Haschick \\& Baan 1985; Koekemoer et al. 1995; Braatz et al. 1994, 1996, 1997; Greenhill et al. 1997a, 2002, 2003; Hagiwara et al. 1997; Falcke et al. 2000). These H$_2$O masers (hereafter `megamasers') appear to be more luminous, by orders of magnitude, then the brightest H$_2$O masers observed in galactic star forming regions. All interferometric data of H$_2$O megamasers point toward a nuclear origin. The emission originates from the innermost parsec(s) of the parent galaxy (e.g. Claussen \\& Lo 1986; Greenhill et al. 1995b, 1996, 1997b; Miyoshi et al. 1995; Claussen et al. 1998; Trotter et al. 1998; Herrnstein et al. 1999). Adopting the so-called unified model (Antonucci \\& Miller 1985; Antonucci 1993) with Seyfert 2 nuclear disks being viewed edge-on, megamaser activity must be related to large column densities along the line of sight, a picture in line with the presence of dense molecular gas.  There exist at least two different classes of megamasers, those forming a circumnuclear disk (like in NGC\\,4258) and those associated with nuclear jets (e.g. NGC\\,1052). The observed H$_2$O transition traces a dense ($\\ga$10$^{7}$\\,cm$^{-3}$) warm ($T_{\\rm kin}$$\\ga$400\\,K) molecular medium. The additional presence of jets and a nuclear X-ray source requires, however, the coexistence of neutral and ionized gas in the nuclear regions of these galaxies. In an attempt to study connections between megamaser emission on milliarcsecond scales and optical phenomena on arcsecond scales, here we present and analyze optical emission line spectra from four megamaser galaxies. \\begin{table*} \\begin{minipage}{180mm} \\caption[]{Journal of the spectroscopic observations} \\begin{flushleft} \\begin{tabular}{lllcllc} \\hline Object/helioc. corr.\\footnote{This heliocentric correction was added to the measured radial velocities.} & Frame & Date (beg.) & p.a. (slit) & Pos. of slit & $\\lambda$ range (\\AA) & Exp.time (s) \\\\ \\hline\\hline NGC\\,1052&F020&11-Jan-96&90\\degr&through nucleus& 5400-7400 & 3600\\\\ $-27$ km/s&F056 & 12-Jan-96 & 90\\degr & through nucleus& 5400-7400 & 1800\\\\ &F057 & 12-Jan-96 & 90\\degr & 2\\arcsec~south of nucleus & 5400-7400 & 1800\\\\ &F058 & 12-Jan-96 & 90\\degr & 2\\arcsec~north of nucleus & 5400-7400 & 1800\\\\ &F078 & 13-Jan-96 & 45\\degr & through nucleus & 5000-6800 & 3600\\\\ &F107 & 14-Jan-96 & 45\\degr & through nucleus & 4000-6000 & 3600\\\\ &F108 & 14-Jan-96 & 45\\degr & 2\\arcsec~northwest of nucleus & 4000-6000 & 3600\\\\ \\hline NGC\\,1368& F023 & 11-Jan-96 & 90\\degr & through nucleus & 5400-7400 & 3600 \\\\ $-17$ km/s & F026 & 11-Jan-96 & 90\\degr & 2\\arcsec~north of nucleus & 5400-7400 & 3600 \\\\ &F080 & 13-Jan-96 & 23\\degr & through nucleus & 5000-6800 & 2700\\\\ &F112 & 14-Jan-96 & 23\\degr & through nucleus & 4000-6000 & 3600\\\\ \\hline Mrk\\,1210 & F066 & 12-Jan-96 & 90\\degr & through nucleus & 5400-7400 & 3600 \\\\ $+6$ km/s &F084 & 13-Jan-96 & 90\\degr & through nucleus & 5000-6800 & 1333\\\\ &F085 & 13-Jan-96 & 90\\degr & 2\\arcsec~south of nucleus & 5000-6800 & 3600\\\\ &F087 & 13-Jan-96 & 90\\degr & 2\\arcsec~north of nucleus & 5000-6800 & 3600\\\\ &F116 & 14-Jan-96 & 90\\degr & through nucleus & 4000-6000 & 3600\\\\ \\hline IC\\,2560 & F030 & 11-Jan-96 & 90\\degr & through nucleus & 5400-7400 & 3600 \\\\ $+20$ km/s & F118 & 14-Jan-96 & 90\\degr & through nucleus & 4000-6000 & 3600 \\\\ \\hline \\end{tabular} \\end{flushleft} \\end{minipage} \\end{table*} ",
        "conclusions": "We have analyzed exploratory optical spectra of four H$_2$O megamaser galaxies in order to disentangle kinematically discernable bulk components of their emission-line regions. For each object we propose a model of the global kinematics. The gas dynamics in IC\\,2560 and NGC\\,1386 is mainly determined by gravitationally dominated motions in a disk. However, there is a bias due to the close-to-edge-on view onto their disks making minor-axis outflow hard to detect. For NGC\\,1052 and Mrk\\,1210 emission from gravity-dominated dynamics of orbiting clouds could be separated from that of systems being in outflow. Line-broadening within the individual bulk components appears to be characterized by scattering wings.  IC\\,2560 is a highly inclined Sb spiral with a nuclear high-ionization Seyfert-2 spectrum displaying relatively narrow (FWHM $\\sim 200$\\,km/s on an arcsecond scale) emission lines from a central rotating disk. If its major-axis orientation and inclination agree with that of the outer galaxy, the core mass (radius 100 pc) will be $9\\,10^6$\\,M$_{\\odot}$, consistently larger than the $2.8\\,10^6$\\,M$_{\\odot}$, given by the orbital velocities of the maser sources presumably embedded in a compact disk of outer radius 0.26\\,pc as measured by Ishihara et al.\\,(2001). A slight blueward asymmetry of the line profiles suggests the presence of outflow (Sect.\\,4.4.1, Fig.\\,3).  The nuclear mass of $\\sim 10^7$\\,M$_{\\odot}$ in IC\\,2560 appears to be rather small. It would be larger by a factor of 1.5 if a distance of 39 Mpc (see comment in Table\\,2) were adopted.  NGC\\,1386 is a highly inclined Sa spiral with a nuclear high-ionization Seyfert-2 spectrum displaying rather structured asymmetric emission-line profiles that we interpret as arising from various parts of a central edge-on orientated disk. We dismiss the term `radiation cone' for this object because the elongated H$\\alpha$ morphology may simply reflect the edge-on view on the matter-bounded disk. The morphology together with analogy reasoning suggests that the disk is a dusty gaseous mini-spiral, which appears to be warped. Its rotation curve inside a diameter of 1.6 kpc yields a dynamically acting mass of $\\sim 5\\,10^9$\\,M$_{\\odot}$, while the visible line-emitting ionized gas has a mass of $\\sim 10^{5-6}$\\,M$_{\\odot}$, which is a very small fraction of the total mass of nuclear gas that is likely to be of the order of ten percent of the dynamical mass.  There is a lack of line emission from the kinematical center (the center of the disk), which is probably due to obscuration by an equatorial dust disk that is evidenced by a dust lane crossing the center. A slight east-west gradient in velocity arising in protrusions seen on an HST image (Ferruit et al. 2000) could be due to minor-axis outflow, i.e. a flow essentially perpendicular to the disk (Sect.\\,4.2.2).  AGN narrow-line emitting disks like in IC\\,2560 and NGC\\,1386 explain general findings that the widths of narrow-lines are broadly correlated with the nuclear stellar velocity dispersion (Nelson \\& Whittle 1996). Such NLRs are likely to be part of the nuclear gas-dust spirals commonly seen on HST images of AGN on a scale of $\\sim10^{3}$ pc (Malkan et al. 1998; Martini \\& Pogge 1999; Martini 2001). However, so far unobserved {\\em very narrow} lines should be emitted from galaxies with central disks seen face-on, if NLRs were {\\em only} dominated by disk dynamics. The other two (non-edge-on) galaxies provide examples for additional dynamical components.  NGC\\,1052 is an elliptical galaxy with a nuclear low-ionization emission line region (LINER). All major line profiles from the brightness center out to 2\\arcsec~northeast can be decomposed into two components $\\cal A$ and $\\cal B$, each with the same width in velocity space for all lines at a given spatial position. $\\cal A$ is attributed to a rotating disk, while $\\cal B$ is blueshifted by $\\sim 400$\\,km/s suggesting outflow of high-density gas relative to the disk. Line ratios of both $\\cal A$ and $\\cal B$ are LINER-like. We propose to relate $\\cal B$ to the peculiar forbidden-line polarized component detected by Barth et al. (1999) (Sect.\\,\\ref{pol}). A faint broad H$\\alpha$ component, which is seen in polarized flux, is not required in total flux. The presence of strong high-density outflow suggests to explore the possibility whether such a flow could be an alternate trigger of the H$_2$O megamaser activity in NGC\\,1052.  Mrk\\,1210 is an amorphously looking Seyfert-2 galaxy that contains a nuclear kiloparsec-spiral seen face-on. A bright emission-line spot appears to be slightly shifted from the partially obscured kinematical center. In total (unpolarized) flux the emission from the bright spot can be decomposed into three narrow-line components $\\cal M$ (the `normal' main component), $\\cal R$ (a redshifted component), and $\\cal B$ (a blueshifted component) without any BLR component. $\\cal M$ is tentatively attributed to the rotating spiral from which only a small velocity variation could be measured because of its near-face-on orientation. $\\cal R$ and $\\cal B$ are outflow components that may both be related to the cospatial starburst and to ionization by the active nucleus. The megamaser activity might be related to the outflow.  The line ratios from the bulk components of Mrk\\,1210 fall into the AGN regime of the VO diagrams. The shifts of component $\\cal{R}$ (by 240\\,km/s to the red of $\\cal{M}$) and of component $\\cal{B}$ (by 169\\,km/s to the blue relative to $\\cal{M}$) are moderate, but of significant influence because major portions of the emission-line region are involved. Regarding line widths and densities $\\cal{B}$ shows characteristics of an intermediate-line-region between BLR and NLR (Sect.\\,\\ref{kin}).  Finally, it should be noted that the successful line-profile decompositions rest on Lorentzians rather than Gaussians as basic functions. Hence, extended wings of {\\em intrinsic narrow-line profiles of the bulk components} may not be uncommon. Crude estimates suggest that electron scattering at the ionized gas itself might lead to a viable explanation, but detailed modelling of such processes would be useful.  There exists a possible connection between optical and radio lines. It is suggested that nuclear outflows impinging onto dense molecular clouds may provide a suitable trigger for megamaser emission. If inner torus and outer disk are misaligned, such an interaction appears to be particularly likely. A convincing verification of such a scenario requires, however, optical spectroscopy with subarcsecond resolution. \\appendix"
    },
    "0212/astro-ph0212173_arXiv.txt": {
        "abstract": "{  We report on a deep XMM-Newton EPIC observation of the globular cluster Omega Cen performed on August 13th, 2001. We have detected 11 and 27 faint X-ray sources in the core and half mass radii, searching down to a luminosity of 1.3 $\\times 10^{31}$ \\ergs\\ in the 0.5-5 keV range. Most sources have bolometric X-ray luminosities between $\\sim 10^{31}-10^{32}$\\ergs. We present the color-color and hardness-intensity diagrams of the source sample, as well as high-quality EPIC spectra of the brightest objects of the field; including the two candidate Cataclysmic Variables (CVs) in the core and the quiescent neutron star low-mass X-ray binary candidate. The spectra of the latter objects fully support their previous classification. We  show that the bulk of sources are hard and spectrally similar to CVs. The lack of soft faint sources might be related to the absence of millisecond pulsars in the cluster. The  XMM-Newton observations reveal the presence of an excess of sources well outside the core of the cluster where several RS CVn binaries have already been found.  We have also analyzed a publicly available Chandra ACIS-I observation performed on January 24-25th, 2000, to improve the \\xmm\\ source positions and to search for source intensity variations between the two data sets.  63 \\xmm\\ sources have a Chandra counterpart, and 15 sources within the half-mass radius have shown time variability. Overall, the general properties of the faint X-ray sources in \\omecen\\ suggest that they are predominantly CVs and active binaries (RS CVn or BY Dra). ",
        "introduction": "Omega Centauri (\\object{NGC 5139}, \\omecen) is one of the best studied objects of our galaxy. It is the most massive globular cluster \\citep[5.1 $\\times 10^{6}$ M$_{\\sun}$, ][]{mey95}. It is characterized by large core and half mass radii \\citep[154.88 $\\arcsec$ and 250.8 $\\arcsec$ respectively,][]{har96}. Binaries are expected to be present in \\omecen\\, either as a result of the evolution of primordial binaries, or through close encounters between stars in the cluster \\citep{dis92,dis94,dav95,ver02}. Binaries such as CVs, low mass X-ray binaries either with a neutron star or a black hole, or active X-ray binaries (RS CVn or BY Dra systems), millisecond pulsars could thus form. Some of these binaries have already been found as faint X-ray sources in \\omecen\\ \\citep{ver01}.  Faint X-ray sources were first detected in \\omecen\\ by the EINSTEIN X-ray satellite. EINSTEIN detected 5 faint point sources \\citep[one in the core,][]{her83} and a possible extended emission region, within the half mass radius \\citep{har82}.  A decade after EINSTEIN, ROSAT detected 22 faint sources in the line of sight of $\\omega$ Cen \\citep{joh94,ver00}. ROSAT confirmed the EINSTEIN sources, and resolved the core source into three components \\citep{ver00}. However, ROSAT did not find any evidence for the diffuse emission seen by EINSTEIN \\citep{joh94}.  More recently, Chandra observed $\\omega$ Cen and detected over 140 faint X-ray sources \\citep{coo02}. From follow-up observations using the accurate Chandra positions, \\citet{coo02} claimed that there were at least three classes of binaries present in the detected sample. Two of the three ROSAT core sources (ROSAT R9a and R9b) may be CVs \\citep{car00}. The third ROSAT core source \\citep[R20,][]{ver00} was associated with a main sequence optical counterpart, showing weak H$\\alpha$ emission, suggesting a BY Dra system rather than a CV. In addition, two more Chandra core sources were detected with HST/WFPC2, with properties matching those of RS CVn or active-corona binaries \\citep{coo02}. Finally, \\citet{coo02} found the X-ray counterparts of two variable binaries discovered far out from the cluster center by \\citet{kal96}. Based on their light curve properties, these two systems are proposed to be RS CVn binaries  \\citep{kal96,kal02}. Beside these three classes of binaries, by looking at the spectral characteristics of the Chandra sources, \\citet{rut02} noticed that one relatively bright object had an extremely soft X-ray spectrum. This spectrum was found to be consistent with those observed from field quiescent neutron star binaries \\citep{rut02}. X-rays would then come from the neutron star surface maintained at a high temperature by episodic mass accretion from a binary companion. In total, up to four different types of binaries may have already been found in the cluster.  We have initiated a survey of nearby globular clusters with XMM-Newton \\citep[M22,][ M13, NGC 6366, Gendre et al., in preparation and $\\omega$ Cen]{web02a}. Because of the limited angular resolution of XMM-Newton, we have selected nearby clusters with large core radii. Taking advantage of the large collecting area of XMM-Newton ($\\sim$ 6 times that of Chandra), we wish to obtain the best possible spectral and timing information for the widest possible sample of faint X-ray sources.  In this paper, we present the first results of our deep \\xmm\\ observation of \\omecen. We describe the general properties of the population of faint X-ray sources detected in the cluster (section \\ref{prop_gen}). Using the publicly available Chandra  observation, we have correlated the \\xmm\\ and Chandra data (sections \\ref{xmm_chandra} and \\ref{diffus}) to improve the \\xmm\\ positions and search for intensity variations between the two data sets (section \\ref{vari}). We also present the spectra of the brightest objects in the field, with the emphasis on those for which an identification already exists (section \\ref{xmm_spec}). We briefly discuss the implications of our findings in section \\ref{discu}. ",
        "conclusions": "\\label{conclu} The main results of our XMM-Newton observation are that in Omega Cen the binaries do not seem be confined to the core or even within the half mass radius, as previously thought, and that the majority of the faint X-ray sources may be CVs, RS CVn or BY Dra binaries. Obviously, the small number of identifications currently available prevents us from reaching definite conclusions. The spectral and timing information provided by XMM-Newton, together with the accurate positions provided by Chandra for a wide sample of faint X-ray sources should encourage follow-up investigations at other wavelengths (optical, UV, radio). In particular, these observations should tell us whether those sources found well outside the core of \\omecen\\ really belong to the cluster. More identifications are also required before reliable comparisons between \\omecen\\ and other clusters with different structural parameters can be drawn.  Such comparisons are critical to a better understanding of the dynamical evolution of globular clusters in general."
    },
    "0212/astro-ph0212459_arXiv.txt": {
        "abstract": "The long-term evolution of massive black hole binaries at the centers of galaxies is studied in a variety of physical regimes, with the aim of resolving the ``final parsec problem,'' i.e. how black hole binaries manage to shrink to separations at which emission of gravity waves becomes efficient. A binary ejects stars by the gravitational slingshot and carves out a loss cone in the host galaxy. Continued decay of the binary requires a refilling of the loss cone. We show that the standard treatment of loss cone refilling, derived for collisionally relaxed systems like globular clusters, can substantially underestimate the refilling rates in galactic nuclei. We derive expressions for non-equilibrium loss-cone dynamics and calculate time scales for the decay of massive black hole binaries following galaxy mergers, obtaining significantly higher decay rates than heretofore. Even in the absence of two-body relaxation, decay of binaries can persist due to repeated ejection of stars returning to the nucleus on eccentric orbits. We show that this recycling of stars leads to a gradual, approximately logarithmic dependence of the binary binding energy on time. We derive an expression for the loss cone refilling induced by the Brownian motion of a black hole binary. We also show that numerical $N$-body experiments are not well suited to probe these mechanisms over long times due to spurious relaxation. ",
        "introduction": "Larger galaxies grow through the agglomeration of smaller galaxies and protogalactic fragments. If more than one of the fragments contains a massive black hole (MBH), the MBHs will form a bound system in the merger product \\citep{Begelman:80,Roos:81,Valtaoja:89}. This scenario has received considerable attention because the ultimate coalescence of such a pair would generate an observable outburst of gravity waves \\citep{thb76}. \\citet{Begelman:80} showed that the evolution of a MBH binary can be divided into three distinct phases: 1. As the galaxies merge, the MBHs sink toward the center of the new galaxy via dynamical friction \\citep{cha43} where they form a binary. 2. The binary continues to decay via gravitational slingshot interactions \\citep{sva74} in which stars on orbits intersecting the binary's are ejected at velocities comparable to the binary's orbital velocity, while the binary's binding energy increases. 3. If the binary's separation decreases to the point where the emission of gravity waves becomes efficient at carrying away the last remaining angular momentum, the MBHs coalesce rapidly.  In this paper we explore Phase 2 and its transition to Phase 3; this transition is understood to be the bottleneck of a MBH binary's path to coalescence.  The picture of \\citet{Begelman:80} has remained essentially valid in the face of several important developments in intervening years.  Early theoretical studies envisioned the centers of galaxies as constant-density cores, with small, embedded stellar cusps that develop around central MBHs via collisional relaxation \\citep{pee72,baw76,Young:77,Ipser:78}. Observations at the time lacked the resolution to verify this picture. In the past decade, space-based observations have revealed that, in the majority of early-type galaxies, the central luminosity density increases continually inward as an approximate power law $\\rho(r)\\sim r^{-\\gamma}$ down to regions where the MBH dominates the gravitational force \\citep{fer94,geb96,res01}. The typical luminosity profile of a faint elliptical galaxy is a nearly-featureless power law with $1.5\\le\\gamma\\le2.5$ at the resolution limit. The times required to build cusps via two-body relaxation are longer than the age of the Universe in most nuclei (e.g.~\\citealt{fab97}) and so relaxation is not a viable explanation for how the cusps formed. In this work we assume that galaxies inherited their steep central density profiles from an early epoch in which the MBH grew.  Cold dark matter (CDM) halos produced by hierarchical structure formation that presently host galaxies and MBHs were cuspy at the time of virialization (e.g.~\\citealt{Ghigna:00,Power:03}). On smaller scales, the stellar cusps may have resulted from gaseous dissipation of star-forming medium inside pre-existing CDM cusps, from adiabatic response of the stars to growth of the MBH, or both. The same processes that led to cusp formation were likely to have fueled the early growth of MBHs and, thus, set the MBHs and galaxies on a course of co-evolution.  We therefore assume that the loci of MBHs have coincided with the centers of cosmic overdensities (dark matter and stellar cusps) since the time MBHs first ascended into the massive range, $M_\\bullet\\sim10^6M_\\odot$.  The cusps that accompany the MBHs endow them with an enlarged effective dynamical mass. \\citet{Milosavljevic:01} found that the dynamical friction time scale (Phase 1) is a function of this effective mass and much shorter than the orbital decay time implied by the masses of the MBHs alone. This is of critical importance for the capacity of MBHs originating in different, merging galaxies to capture each other and form a hard binary in the resulting galaxy. A MBH binary is customarily called ``hard'' if its binding energy per unit mass exceeds the specific kinetic energy of ambient stars, $G\\mu/4a\\gtrsim\\sigma^2$, where $\\mu=M_1M_2/(M_1+M_2)$ is the reduced mass, $a$ is the binary's semi-major axis and $\\sigma$ is the 1D stellar velocity dispersion. \\citet{Milosavljevic:01} also found that in mergers of equal-mass $\\rho(r)\\sim r^{-2}$ galaxies, the true separations at which the binaries settle at the beginning of Phase 2 are $\\sim 10$ times smaller than implied by arguments based on the equipartition of energy.  Masses of MBH's correlate exceptionally well with the central line-of-sight velocity dispersions of their host galaxies \\citep{fem00,geb00}. It is likely that the masses are also related to the depth of the potential wells associated with the dark matter halos in which the host galaxies are embedded \\citep{Ferrarese:02a}, which justifies attempts to characterize the {\\it evolution} of the MBH population using the statistics of hierarchical formation of structure in dark matter cosmologies \\citep{kah00,Menou:01,Volonteri:02}.  As a corollary, one can estimate the frequency of MBH binary coalescence at redshifts that are beyond the reach of existing telescopes but within reach of a future space-based gravity wave detector such as LISA.\\footnote{Laser Interferometer Space Antenna, http://lisa.jpl.nasa.gov} The coalescence of two $10^5-10^6 M_\\odot$ MBHs occurring practically anywhere in the observable Universe would generate a significant signal for LISA.  Early calculations (e.g.~\\citealt{Haehnelt:94}) identified the MBH coalescence event rate with the galaxy merger rate.  It was subsequently noted that stellar-dynamical processes are inefficient at facilitating the decay of MBH binaries when their separations shrink below about one parsec and that the ultimate coalescence should not be taken for granted \\citep{val96,gor00,Milosavljevic:01,Yu:02a}. Circumstantial evidence, however, suggests that long-lived MBH binaries are in fact rare. The coincidence of mean black hole mass densities derived from kinematical studies of local galaxies and distant quasars (\\citet{Ferrarese:02b}) suggests that MBHs are rarely ejected from galaxy centers by the strong interactions that would accompany the infall of a third MBH into a nucleus containing an uncoalesced binary. Furthermore there is evidence for MBH coalescences in the morphology of radio galaxies, at a rate approximately equal to the galaxy merger rate \\citep{MerrittEkers:02}.  There are many processes that could drive the MBH binaries that form in galaxy mergers to coalescence. In this paper we present some new ideas regarding one such process, the gravitational slingshot ejection of stars.  Ejection of stars from the galaxy nucleus will result in a depletion of stars in the ``loss cone.'' We define the loss cone as the phase space domain in which orbits approach the binary closely enough to be influenced by the near field of the individual binary components. The orbits in the loss cone are subject to gravitational capture and slingshot ejection.  The loss cone has thus traditionally been viewed as a sink in the phase space, surrounded by a distribution of stars in equilibrium with respect to two-body gravitational encounters.  The latter assumption allows a steady-state solution for the phase space density near the loss cone boundary to be derived (e.g.~\\citealt{cok78}); this steady state reflects a balance between the supply of stars to the loss cone by gravitational scattering and the loss due to capture or ejection. In this standard picture, stars that scatter into the loss cone are ejected from the system with net energy loss, resulting in a more tightly bound binary. After all of the stars originally inside the loss cone have been removed, continued decay of the binary hinges on the flux of stars scattering into the loss cone.  In this paper we study the physics of the ``worst case scenario,'' involving nearly spherical galaxies, in which the MBH binaries are least likely to coalesce in their lifetime. In \\S~\\ref{sec:reviewck} we review the standard solution for the flux. In \\S~\\ref{sec:issues} we question the applicability of some of the assumptions that were made to derive that solution. In \\S~\\ref{sec:timedep} and \\S~\\ref{sec:secondary} we introduce some novel aspects of the dynamics of stars in the loss cone and outline an iterative procedure for calculating the long-term evolution of a massive MBH binary in a spherical galaxy.  The details emphasized here tend to shorten the timescale for the decay of MBH binaries, sometimes significantly. In \\S~\\ref{sec:nbody} we evaluate the effectiveness of $N$-body simulations to model the long-term evolution of MBH binaries. In \\S~\\ref{sec:brownian} we address the implications of the Brownian motion of MBH binaries for the long-term evolution.  In \\S~ \\ref{sec:coalescence} we review the likelihood for the coalescence of MBH binaries. Extensions to axisymmetric and triaxial galaxies, as well as applications to specific galaxy models, are deferred for future work. ",
        "conclusions": "We conclude with a tabulated summary of physical regimes that have been discussed so far (Table \\ref{tab:regimes}).  \\begin{deluxetable}{lll} \\tablecolumns{3} \\tablewidth{30pc} \\tablehead{ \\colhead{} & \\colhead{Form of Decay} & \\colhead{Regime}} \\tablecaption{Physical Regimes for Long-Term Decay of Massive Black Hole Binaries} \\startdata $a^{-1} \\propto$ & $t + \\textrm{const}$ & pinhole \\\\ $ $ & $t/N^\\alpha + \\textrm{const}$,\\ \\ \\ \\ ($0<\\alpha<1$) & pinhole \\& diffusion \\\\ $ $ & $t/N$ & diffusion \\\\ $ $ & $\\textrm{const}$ & large-$N$ limit \\\\ $ $ & $\\ln(1 + t/t_0) + \\textrm{const}$ & pure reejection \\\\ $ $ & $\\sqrt{1+Dt}$ & pinhole with reejection \\enddata \\label{tab:regimes} \\end{deluxetable}  Clearly the long-term evolution of a massive black hole binary can be very different in different environments. We distinguish three characteristic regimes.  1. {\\it Collisional.} The relaxation time is shorter than the lifetime of the system and the phase-space gradients at the edge of the loss cone are given by steady-state solutions to the Fokker-Planck equation (\\S~\\ref{sec:reviewck}). The densest galactic nuclei may be in this regime. Resupply of the loss cone takes place on the time scale associated with scattering of stars onto eccentric orbits. Most of this scattering occurs in the ``diffusive'' (as opposed to ``pinhole'') regime and the decay time of a MBH binary scales as $|a/\\dot a|\\sim m_*^{-1}\\sim N$. In the densest galactic nuclei, collisional loss cone refilling may just be able to drive a MBH binary to coalescence in a Hubble time (e.g.~\\citealt{Yu:02a}). $N$-body studies with any feasible $N$ are guaranteed to be in the collisional regime, and may also be in the pinhole regime corresponding to an always-full loss cone, leading to artifically high decay rates, $a^{-1}\\sim t$ (\\S~\\ref{sec:nbody}). In $N$-body studies with $N\\lesssim 10^6$, Brownian motion of the binary further enhances the decay (\\S~\\ref{sec:brownian}).  2. {\\it Collisionless.} The relaxation time is longer than the system lifetime and gravitational encounters between stars can be ignored. Examples are the low-density nuclei of bright elliptical galaxies; dark matter density cusps are also in this regime if the dark matter particles are much less massive than stars. The MBH binary quickly interacts with stars whose pericenters lie within its sphere of influence; in a low-density nucleus, the associated mass is less than that of the MBH binary and the decay tends to stall at a separation too large for gravity wave emission to be effective. However evolution can continue due to reejection of stars that lie within the binary's loss cone but have not yet escaped from the system. In the spherical geometry, reejection implies $|a/\\dot a|\\sim (1+t/t_0)/a$, leading to a logarithmic dependence of binary hardness on time (\\S~\\ref{sec:secondary}). Reejection in galactic nuclei may contribute a factor of $\\sim$ a few to the change in $a$ in MBH binaries over a Hubble time.  3. {\\it Intermediate.} The relaxation time is of order the age of the system. While gravitational encounters contribute to the re-population of the loss cone, not enough time has elapsed for the phase space distribution to have reached a collisional steady state. We argued that most galactic nuclei are likely to be in this regime. The flux of stars into the loss cone can be substantially higher than predicted by the steady-state theory, due to strong gradients in the phase space density near the loss cone boundary produced when the MBH initially formed (\\S~\\ref{sec:timedep}). This transitory enhancement would be most important in a nucleus that continues to experience mergers or infall, in such a way that the loss cone repeatedly returns to an unrelaxed state with its associated steep gradients.  The evolution of a real MBH binary may reflect a combination of the different regimes summarized above, as well as other mechanisms that we have discussed only briefly or not at all. We note a close parallel between the ``final parsec problem'' and the problem of quasar fueling: both requre that of order $10^8M_\\odot$ be supplied to the inner parsec of a galaxy in a time shorter than the age of the universe. Nature clearly accomplishes this in the case of quasars, probably through gas flows driven by torques from stellar bars \\citep{Begelman:90}. The same inflow of gas could contribute to the decay of a MBH binary in a number of ways: by leading to the renewed formation of stars which subsequently interact with the binary; by inducing torques which extract angular momentum from the binary \\citep{Armitage:02}; through accretion, increasing the masses of one or both of the MBHs and reducing their separation; etc.  In massive elliptical galaxies or spheroids, the fraction of mass in the form of gas during the last major merger event is likely to have been small \\citep{kah00}. In addition to the mechanisms discussed in this paper, decay of a MBH binary could be enhanced by the presence of a significant population of ``centrophilic'' orbits, orbits that pass near the center of the potential once per crossing time. Such orbits may constitute a large fraction of the mass of a triaxial nucleus \\citep{Poon:02}. Even in galaxies where none of these mechanisms is effective and the decay stalls, infall of a third MBH following a merger or accretion event could accelerate the decay via the Kozai mechanism \\citep{Blaes:02}, or the strong three-body gravitational interactions \\citep{Makino:94,Valtonen:96}, which induce large changes in the binary's eccentricity and enhanced rates of gravity wave emission at pericenter.  We have not attempted in this paper to make quantitative estimates of the decay rates or stalling radii of MBH binaries in real galaxies. Such a program would be problematic for a number of reasons, most importantly the uncertain history of the loss cone and the unknown elapsed time and character of the most recent merger (\\S~\\ref{sec:timedep}). However we argued that a recent study \\citep{Yu:02a} could have overestimated stalling radii, due both to the neglect of the various new mechanisms discussed here, and also by ignoring the influence that MBH formation must have had on the form of the nuclear density profile: this  profile must have been steeper before the MBH binary formed, leading to substantially enhanced loss cone fluxes and more rapid decay early in the life of the binary. Nevertheless, none of our results would lead us to rule out the existence of uncoalesced MBH binaries in at least some galaxies, particularly galaxies with large, low-density nuclei and a little gas (\\S~\\ref{sec:coalescence}).  A striking and disappointing conclusion of this study is the difficulty of using numerical $N$-body simulations to follow the long-term evolution of MBH binaries (\\S~\\ref{sec:nbody}). We argued that a variety of collisional effects associated with small $N$ would lead to a more rapid decay of the binary than in real galaxies. While there is still a useful role for $N$-body simulations in following the initial formation and early evolution of a MBH binary, when much of the decay and associated cusp destruction takes place, the rapid decay of $a(t)$ seen in most published simulations is due to loss-cone repopulation occuring at rates much higher than expected in real galaxies. We believe that the future role of $N$-body simulations in this field will be limited to predicting the form of the phase-space gradients produced during the formation of a MBH binary, which can then be adapted as initial conditions for a collisionless or Fokker-Planck treatment. A carefully designed set of $N$-body experiments might also be used to understand the $N$-dependence of collisional loss cone effects which could then be extrapolated to the larger-$N$ regime characteristic of real galaxies.  \\vspace{0.5cm} We are grateful to A. Kosowsky, C. Pryor, and R. Spurzem for their detailed comments on an early version of the manuscript.  We acknowledge helpful discussions with J. Binney, S. Phinney, F. Rasio, and M. Rees. This work was supported by NSF grants AST 96-17088 and AST 00-71099 and by NASA grants NAG5-6037 and NAG5-9046 and STScI grant HST-AR-08759. M.~M.\\ thanks the Sherman Fairchild Foundation for support."
    },
    "0212/physics0212072_arXiv.txt": {
        "abstract": "Using magnetic field and plasma data acquired during {\\it Phobos-2} mission in regions which have not been explored before, we study the solar wind interaction with Phobos. The draping magnetic field of the solar wind around Phobos appears at distances of 200--300~km from the Phobos day-side due to a density and magnetic field pile up in front of the Phobos obstacle. The nature of the interaction and the magnetic field signatures observed are consistent with the ratio of the proton skin depth to the actual size of the Phobos obstacle to the solar wind. Phobos deflects the flow of the solar wind and the subsolar stand-off distance of the deflection is about 16--17 Phobos radii. Source with equivalent magnetic moment $M'\\simeq10^{15}$ A$\\cdot$m$^2$ in Phobos leads to the development of such an obstacle to solar wind flow around Phobos. These data give a lesson for the study of the interaction of the small, magnetized object with solar wind. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/astro-ph0212109_arXiv.txt": {
        "abstract": "{ We present results of mid-infrared $\\lambda = 5.0 - 16.5~\\mathrm{\\mu m}$ spectrophotometric imaging of the starburst galaxies \\object{M\\,82}, \\object{NGC\\,253}, and \\object{NGC\\,1808} from the ISOCAM instrument on board the {\\em Infrared Space Observatory\\/}.  The mid-infrared spectra of the three galaxies are very similar in terms of features present. The $\\lambda \\ga 11~\\mathrm{\\mu m}$ continuum attributed to very small dust grains (VSGs) exhibits a large spread in intensity relative to the short-wavelength emission.  We find that the 15$\\,\\mu$m dust continuum flux density correlates well with the fine-structure [\\ion{Ar}{ii}] 6.99$\\,\\mu$m line flux and thus provides a good quantitative indicator of the level of star formation activity.  By contrast, the $\\lambda = 5 - 11~\\mathrm{\\mu m}$ region dominated by emission from polycyclic aromatic hydrocarbons (PAHs) has a nearly invariant shape.  Variations in the relative intensities of the PAH features are nevertheless observed, at the $20\\% - 100\\%$ level.  We illustrate extinction effects on the shape of the mid-infrared spectrum of obscured starbursts, emphasizing the differences depending on the applicable extinction law and the consequences for the interpretation of PAH ratios and extinction estimates.  The relative spatial distributions of the PAH, VSG, and [\\ion{Ar}{ii}] 6.99$\\,\\mu$m emission between the three galaxies exhibit remarkable differences.  The $\\la 1~\\mathrm{kpc}$ size of the mid-infrared source is much smaller than the optical extent of our sample galaxies and $70\\% - 100\\%$ of the {\\em IRAS\\/} 12$\\,\\mu$m flux is recovered within the ISOCAM $\\leq 1.5~\\mathrm{arcmin^{2}}$ field of view, indicating that the nuclear starburst dominates the total mid-infrared emission while diffuse light from quiescent disk star formation contributes little. ",
        "introduction": "\\label{Sect-intro}  As hosts of vigorous episodes of star formation activity, starburst galaxies in the local Universe offer a unique opportunity to study the star formation process in extreme environments reminiscent of the conditions presumably prevailing in primordial galaxies.  Studies in recent years have underscored the prominent role of starbursts in galaxy formation and evolution as well as their significant contribution to the extragalactic background, and to the chemical enrichment and heating of the intergalactic medium up to the highest redshifts (e.g. Steidel et al.~\\cite{Ste96}; Puget et al.~\\cite{Pug96}; Smail et al.~\\cite{Sma97}; Franx et al.~\\cite{Fra97}; Pettini et al.~\\cite{Pet98}; Elbaz et al.~\\cite{Elb02}). Understanding the starburst phenomenon is thus a key issue of local and cosmological relevance.  Nuclear starburst galaxies generally emit the bulk of their luminosity in the infrared, primarily as reprocessed radiation by interstellar dust grains heated by important populations of hot massive stars.  In addition, a substantial fraction -- if not most -- of the star formation in starburst systems is heavily obscured at optical and ultraviolet wavelengths by large amounts of dust.  This has been particularly dramatically illustrated by mid-infrared imaging of the colliding pair \\object{NGC\\,4038/4039} (Mirabel et al.~\\cite{Mir98}).  Studies at infrared wavelengths are thus crucial in defining the properties of starburst galaxies and the interplay between the primary energy sources and the reprocessing material.  The growing evidence for the presence and significant role of dust in distant galaxies further emphasizes the need for infrared investigations of nearby dust-rich templates (e.g. Rowan-Robinson \\cite{Row01}; Blain et al.~\\cite{Bla02}; and references therein). However, except for a limited sample, the detailed spatial and spectral energy distributions of starburst galaxies in this regime are still poorly known. Progress has been hindered by instrumental limitations, and by the difficulties inherent to ground-based and air-borne observations at these wavelengths.  The {\\em Infrared Space Observatory\\/} ({\\em ISO\\/}; Kessler et al.~\\cite{Kes96}) has revolutionized the field of infrared astronomy (see the reviews by Cesarsky \\& Sauvage \\cite{Ces99}; Genzel \\& Cesarsky \\cite{Gen00}), with unprecedented capabilities compared to its predecessor the {\\em Infrared Astronomical Satellite} ({\\em IRAS}\\/).  In the mid-infrared (MIR, $\\lambda = 5.0 - 16.5~\\mathrm{\\mu m}$), the camera ISOCAM (Cesarsky et al.~\\cite{Ces96}) provided an improvement in sensitivity and spatial resolution by factors of $\\sim 1000$ and $\\sim 60$, respectively. In addition to broad- and narrow-band filters, ISOCAM featured continuously variable filters (CVFs) allowing arcsecond-scale spectrophotometric imaging with full wavelength coverage at a resolution of $R \\equiv \\lambda/\\Delta\\lambda \\sim 40$.  The MIR regime contains a variety of key features tracing different components of the interstellar medium (ISM) and is subject to relatively little extinction by dust ($A_\\mathrm{5 - 17\\,\\mu m} < 0.1\\,A_{V}$). Spectrophotometric observations permit studies of the most prominent features to characterize the spatial distribution and infer the properties of the ISM and exciting sources, deep into obscured star-forming regions.  In this paper, we present results of ISOCAM CVF observations of \\object{M\\,82}, \\object{NGC\\,253}, and \\object{NGC\\,1808}. \\object{M\\,82} and \\object{NGC\\,253} are two of the nearest and brightest starburst galaxies (at distances of $D = 3.3$ and 2.5~Mpc, respectively, $1^{\\prime\\prime} = 16$ and 12~pc; Freedman \\& Madore \\cite{Fre88}; Davidge \\& Pritchet \\cite{Dav90}). Both have been extensively studied in all accessible spectral regions and have often been considered as prototypical objects for the starburst phenomenon (see Telesco \\cite{Tel88}, Rieke et al.~\\cite{Rie93}, and Engelbracht et al.~\\cite{Eng98} for reviews). \\object{NGC\\,1808} is a more distant system ($D = 10.9~\\mathrm{Mpc}$ for $H_{0} = 75~\\mathrm{km\\,s^{-1}\\,Mpc^{-1}}$, $1^{\\prime\\prime} = 53~\\mathrm{pc}$; Sandage \\& Tammann \\cite{San87}) which closely resembles \\object{M\\,82} and \\object{NGC\\,253} in its global properties (e.g. Dahlem et al.~\\cite{Dah90}; Junkes et al.~\\cite{Jun95}). In all three galaxies, starburst activity takes place within the central $\\sim 0.1 - 1~\\mathrm{kpc}$ regions which are severely obscured at optical and ultraviolet wavelengths.  The large extinction levels likely result in part from high disk inclination angles of $80^{\\circ}$, $78^{\\circ}$, and $57^{\\circ}$ for \\object{M\\,82}, \\object{NGC\\,253}, and \\object{NGC\\,1808} (G\\\"otz et al.~\\cite{Got90}; Pence \\cite{Pen81}; Reif et al.~\\cite{Rei82}).  At comparable infrared luminosities\\footnote{as computed from the {\\em IRAS\\/} fluxes following the prescription of Sanders \\& Mirabel (\\cite{San96}).} $L_\\mathrm{IR} \\equiv L_\\mathrm{8-1000\\,\\mu m}$ of $3.7 \\times 10^{10}$, $1.4 \\times 10^{10}$, and $3.4 \\times 10^{10}~\\mathrm{L_{\\odot}}$ for \\object{M\\,82}, \\object{NGC\\,253}, and \\object{NGC\\,1808}, these galaxies also share key morphological features including a kiloparsec-scale stellar bar, large amounts of molecular gas in the central $\\sim 1~\\mathrm{kpc}$ predominantly concentrated in circumnuclear ring-, spiral-, or bar-like structures, populations of compact and luminous ``super star clusters'' as well as of compact non-thermal radio sources associated with supernova remnants, and a large-scale outflowing starburst wind (F\\\"orster Schreiber et al.~\\cite{FS01}; Engelbracht et al.~\\cite{Eng98}; Tacconi-Garman et al.~\\cite{Tac96}; and references therein). The triggering of starburst activity in \\object{M\\,82} is generally attributed to gravitational interaction as evidenced by the extended filamentary tidal features in \\ion{H}{i} gas distribution threading \\object{M\\,82} and its neighbours \\object{M\\,81} and \\object{NGC\\,3077} or, alternatively, to the stellar bar which may itself have been induced by the interaction (e.g. Yun et al.~\\cite{Yun93, Yun94}). In \\object{NGC\\,253}, the bar itself appears to be the primary mechanism responsible for the onset of starburst activity as there are no obvious signatures of interaction with a nearby companion, although a past merger or accretion event involving a small galaxy has been suggested based on kinematical evidence (e.g. Anantharamaiah \\& Goss \\cite{Ana96}; B\\\"oker et al.~\\cite{Bok98}). Tidal interaction with \\object{NGC\\,1792} has been proposed for \\object{NGC\\,1808} based on circumstantial evidence (e.g. Dahlem et al.~\\cite{Dah90}; Koribalski et al.~\\cite{Kor93}). The recent high sensitivity \\ion{H}{i} observations by Dahlem et al. (\\cite{Dah01}) do not indicate the presence of any tidal feature, dwarf galaxy or other debris that could support a close passage in the past, leaving the stellar bar as more probable trigger or a possible merger/accretion event.  In view of their similar nature, we focus on the comparison of the spatial and spectral properties of \\object{M\\,82}, \\object{NGC\\,253}, and \\object{NGC\\,1808} at MIR wavelengths. Though small, this sample also allows us to probe the transition regime between normal spiral and irregular galaxies and the more extreme luminous and ultraluminous infrared galaxies (LIRGs and ULIRGs for which $L_\\mathrm{IR} > 10^{11}$ and $> 10^{12}~\\mathrm{L_{\\odot}}$, respectively). Section~\\ref{Sect-obs} describes the observations and data reduction procedure, and Sect.~\\ref{Sect-res} presents the results. Section~\\ref{Sect-Spat_distr} discusses the origin and spatial distibution of the continuum and emission features. Section~\\ref{Sect-Spec_distr} addresses the issues of extinction effects, variations in spectral properties, and indicators of star formation activity. Section~\\ref{Sect-conclu} summarizes the paper. ",
        "conclusions": "\\label{Sect-conclu}  We have presented ISOCAM $\\lambda = 5.0 - 16.5~\\mathrm{\\mu m}$ spectrophotometric imaging of the starburst galaxies \\object{M\\,82}, \\object{NGC\\,253}, and \\object{NGC\\,1808}.  The spectrum of all three objects, down to the smallest scales of $\\sim 100~\\mathrm{pc}$ accessible from our data, exhibit similar characteristics including prominent PAH bands, a featureless continuum most obvious at $\\lambda \\ga 11~\\mathrm{\\mu m}$, and a trough in the 10$\\,\\mu$m region. We securely identified the main emission features detected in the $R \\sim 40$ ISOCAM data of \\object{M\\,82} and \\object{NGC\\,253} based on their $R \\sim 500 - 1000$ SWS spectra. The comparison emphasizes the caution that should be exercised when interpreting low resolution data because of potential blends or misidentifications between emission lines from ionized gas and PAH features originating in PDRs, a notable example being [\\ion{Ne}{ii}] 12.81$\\,\\mu$m and PAH 12.7$\\,\\mu$m.  Using a simple model combining a template PDR spectrum and a power-law $f_{\\nu} \\propto (\\lambda - 8.5)^{1.5}$, we constructed a representative starburst SED and explored the effects of extinction at MIR wavelengths.  Our simulations illustrate the importance of the assumed extinction law (e.g. the widely used Draine \\cite{Dra89} law versus the Galactic Center law of Lutz \\cite{Lut99}) and of the intrinsic PAH spectrum (especially the gap between the main $6 - 9~\\mathrm{\\mu m}$ and $11 - 13~\\mathrm{\\mu m}$ complexes) in shaping the SED of astronomical sources.  This complicates the interpretation of PAH ratios as well as extinction measurements relying on the silicate dust absorption at 9.7$\\,\\mu$m (see also Sturm et al.~\\cite{Stu00}).  As observed previously in a wide range of Galactic and extragalactic sources, the $5 - 11~\\mathrm{\\mu m}$ spectrum in our galaxies is nearly invariant. The relative PAH intensities exhibit nevertheless measureable and significant variations of $20\\% - 100\\%$ which may be attributed to various, possibly interrelated effects including the intensity of the incident radiation field and the PAH size distribution, ionization, and dehydrogenation. In our sample, \\object{M\\,82} probably best illustrates variations of PAH ratios due to an increased fraction of ionized PAHs within the most intense starburst sites and, admittedly speculatively, perhaps also to differences in typical PAH sizes depending on the molecular gas concentrations. The PAH~7.7$\\,\\mu$m L/C ratio in all three galaxies clearly lies in the range observed for pure starburst systems and extends the trend reported previously by Rigopoulou et al.~(\\cite{Rig99}) to lower luminosities.  In contrast, the $\\lambda \\ga 11~\\mathrm{\\mu m}$ region varies most among our sample galaxies and the ISOCAM maps show a comparatively more compact 15$\\,\\mu$m continuum distribution relative to the PAH emission. Strong 15$\\,\\mu$m continuum in \\object{M\\,82} and \\object{NGC\\,253} indicates an important contribution by VSGs contrary to the case of \\object{NGC\\,1808} where the long-wavelength emission is much flatter presumably because of negligible contribution by VSGs. We find, however, that in all three galaxies the 15$\\,\\mu$m continuum and [\\ion{Ar}{ii}] 6.99$\\,\\mu$m line fluxes satisfy a linear relationship.  We infer from this that the 15$\\,\\mu$m continuum provides a good indicator of star formation activity in starbursts, complementing the similar results of Roussel et al.~(\\cite{Rou01b}) for galactic disks. In a broader perspective, our galaxies fit well in the trend of increasing ISOCAM 15$\\,\\mu$m/7$\\,\\mu$m ratios with higher levels of star formation activity found among normal disk galaxies and starburst-powered LIRGs/ULIRGs (e.g. Laurent et al.~\\cite{Lau00}; Roussel et al.~\\cite{Rou01b}; Dale et al.~\\cite{Dal01}).  The value of the ISOCAM spectrophotometric imaging of \\object{M\\,82}, \\object{NGC\\,253}, and \\object{NGC\\,1808} presented in this paper also lies in that it complements existing MIR data with maps of PAH features, fine-structure lines, and continuum components not previously imaged. The poorer angular resolution of ISOCAM compared to that achieved with large ground-based telescopes is compensated by the larger field of view revealing more of the large scale emission.  In that respect, we stress the small size of the MIR source relative to the optical extent of all three galaxies: the ISOCAM maps cover the central $\\approx 1.5~\\mathrm{kpc}$ for \\object{M\\,82}, 0.6~kpc for \\object{NGC\\,253}, and 5~kpc for \\object{NGC\\,1808} while the optical diameters are about 10~kpc for \\object{M\\,82} and 20~kpc for the other two. By measuring the flux density within the ISOCAM LW10 filter bandpass, equivalent to the {\\em IRAS\\/} 12$\\,\\mu$m band, we recover all of the {\\em IRAS\\/} 12$\\,\\mu$m emission in the entire field of view for \\object{M\\,82} and \\object{NGC\\,1808}, and about 70\\% for \\object{NGC\\,253}. This suggests that the total MIR emission is strongly dominated by the starburst sites in the nuclear regions while the more quiescent star formation taking place in the disk at larger radii does not contribute much. It will be interesting to see whether the MIR cameras on board SIRTF, to be launched in 2003, will confirm this result or discover faint diffuse emission, especially at shorter ($3 - 5~\\mathrm{\\mu m}$) wavelengths."
    },
    "0212/astro-ph0212423_arXiv.txt": {
        "abstract": "We present multi-frequency VLBA\\footnote[2]{The Very Long Baseline Array (VLBA) and Very Large Array (VLA) are operated by the National Radio Astronomy Observatory, a facility of the National Science Foundation operated under cooperative agreement by Associated Universities for Research in Astromomy, Inc.} observations of PMN~J1632--0033, one of the few gravitationally lensed quasars suspected of having a central ``odd'' image.  The central component has a different spectral index than the two bright quasar images. Therefore, either the central component is not a third image, and is instead the active nucleus of the lens galaxy, or else it is a third image whose spectrum is inverted by free-free absorption in the lens galaxy.  In either case, we have more constraints on mass models than are usually available for a two-image lens, especially when combined with the observed orientations of the radio jets of the two bright quasars.  If there is no third quasar image, the simplest permitted model is a singular isothermal sphere in an external shear field: $\\beta= 2.05^{+0.23}_{-0.10}$ ($2\\sigma$), where $\\rho(r) \\propto r^{-\\beta}$.  If the central component is a third image, a hypothesis which can be tested with future high-frequency observations, then the density distribution is only slightly shallower than isothermal: $\\beta=1.91\\pm 0.02$ ($2\\sigma$).  We also derive limits on the size of a constant-density core, and the break radius and exponent of an inner density cusp. ",
        "introduction": "\\label{sec:intro}  The central density profile $\\rho(r)$ of galaxies is poorly known, but is important for understanding galaxy structure and the nature of dark matter.  Theoretical cold-dark-matter (CDM) models predict a steep central density cusp, $\\rho \\propto r^{-\\beta}$ with $\\beta \\simeq$~1--1.5 \\citep[see, e.g.,][]{nfw97,moore99}.  However, measurements of spiral galaxy rotation curves suggest the dark matter has a constant-density core.  This has been argued most strongly for low-surface-brightness galaxies \\citep{deblok01} but may hold for all spiral galaxies \\citep{ds00,sb00}, including our own Galaxy \\citep{be01}.  The conflict may be due to observational effects such as beam smearing, at least in part \\citep{vs01}, but the apparent conflict has encouraged some theorists to invent mechanisms that avoid producing cusps (e.g., bar-induced evolution, Weinberg \\& Katz 2002; warm dark matter, Bode, Ostriker, \\& Turok 2001; self-interacting dark matter, Spergel \\& Steinhardt 2000).  For elliptical galaxies, where rotation curves cannot be measured, the central density profile is even less well understood.  Observations of early-type galaxies with the Hubble Space Telescope (HST) show that the light distribution is cuspy, with cusp parameters that may be correlated with global galaxy properties \\citep{lauer95,faber97}.  But such studies do not probe mass directly, and are limited to nearby galaxies due to the requirements on angular resolution.  Multiple-image gravitational lenses can be used to measure the central surface density of galaxies at significant redshift, by searching for faint central images of the background object.  In most cases, lenses should produce an odd number of images \\citep{dr80,burke81}, one of which is faint and appears near the center of the lens galaxy.  This ``central image'' or ``odd image'' corresponds to the central maximum in the time-delay surface.  However, of the $\\sim$80 known systems in which a galaxy produces multiple images of a quasar or radio source, almost all consist of either 2 or 4 images.  One system, APM~08279+5255, definitely has three quasar images \\citep{lewis02b}, but the third image may be due to a ``naked cusp'' configuration rather than a central time-delay maximum \\citep{lewis02a}.  No central image has been identified definitively, even among the radio lenses, where the central image could in principle be seen through the light and dust of the lens galaxy.  The absence of central images has been attributed to the demagnification of the image by the high central surface density of the lens galaxy \\citep[see, e.g.,][]{nsc86,wn93,rm01,eh02,keeton02}.  An ``extra'' radio component has been detected in several radio lenses, and in three of those cases it is still unclear whether the extra component is an additional quasar image or faint radio emission from an active galactic nucleus (AGN).  Those three systems are 0957+561 \\citep{harvanek97}, MG~J1131+0456 \\citep{chen93}, and the subject of this paper: PMN~J1632--0033 \\citep{winn02}.  Radio components A and B of J1632--0033 are two images of a $z=3.42$ quasar, but the third, faint component C is of unknown origin (see Fig~\\ref{fig:merlin}).  In an HST image, the lens galaxy appears to be a fairly circular early-type galaxy with an optical effective radius of $\\approx 0\\farcs2$, but it is too faint to be characterized in detail.  The position of component C is near the optical lens galaxy position, as expected for either an AGN or a central odd image.  If A, B, and C are all images of the same quasar, then they should have similar radio spectra.  However, \\citet{winn02} only presented measurements of the flux density of C at one frequency.  \\begin{figure}[h] \\figurenum{1} \\epsscale{0.5} \\plotone{f1.eps} \\caption{ Radio map of \\lens reproduced from \\citet{winn02}, illustrating the configuration of components A, B, and C.  The map is based on the 5.0~GHz {\\sc merlin} observation of 2001~April~13.  The restoring beam ($88\\times 44$~mas, P.A.\\ $24\\arcdeg$) is illustrated in the upper left corner.  Contours begin at $3\\sigma$ and increase by factors of two, where $\\sigma=0.13$~mJy~beam$^{-1}$. } \\label{fig:merlin} \\end{figure}  Here we present new VLBA data that, in combination with the previous data, allow us to measure the flux densities of all three components at four widely spaced frequencies.  Our analysis shows that C has a significantly different radio spectrum than A and B.  We argue that there are two alternative explanations for this difference.  Either C is not a third quasar image, and instead is the active nucleus of the lens galaxy; or, C is a third image, but its spectral index is inverted by free-free absorption in the lens galaxy.  Whichever of these possibilities is true, the additional information from the VLBA data provides more constraints on mass models of the lens galaxy than are usually available for a two-image lens.  If the central component is a third image, then its position and flux are valuable constraints.  If the central component is the lens galaxy, then our data provide an unusually precise position for the lens galaxy, and a strong upper limit on the flux of a third image.  The VLBA maps also reveal the relative orientations of the radio jets associated with A and B.  As with all gravitational lensing studies, we still do not have enough information to determine the complete density profile of the lens galaxy, but we can use parameterized models to explore the implications of the data.  Historically, two classes of parameterized models have been used to study the central-image problem: models with constant-density cores, and models with central cusps.  Early studies used softened power laws, $\\rho \\propto (r^2+r_c^2)^{-\\beta/2}$, and found that central images are suppressed if the core radius $r_c$ is sufficiently small \\citep[see, e.g.,][]{wn93}.  Since the demagnification of the central image is determined by the central surface density, steeper density profiles (larger $\\beta$) allow for larger core radii for the same flux limit on the central image.  Power laws shallower than isothermal ($\\beta < 2$) produce central images even in the limit $r_c \\rightarrow 0$.  With the gradual realization that theoretical dark-matter profiles and observed surface-brightness profiles of early-type galaxies have cusps rather than cores, attention has shifted to the properties of density distributions with central cusps.  \\citet{rm01} used the absence of odd images in the {\\sc class} sample, presently the largest homogeneous sample of radio lenses, to argue that the global mass profiles of the lens galaxies cannot be much shallower than isothermal.  \\citet{mkk01} reached a similar conclusion about the particular radio lens {\\sc class}~B1933+503.  \\citet{keeton02} found that the distribution of stars observed in HST images of nearby early-type galaxies is often sufficiently concentrated to suppress the central image without recourse to dark matter.  We describe our observations of \\lens in \\S~\\ref{sec:observations}, and our interpretation of the results in \\S~\\ref{sec:interpretation}. In \\S~\\ref{sec:models}, we compute lens models for this system under the two different hypotheses for the nature of the central component. We consider density distributions that are scale-free power laws, and also determine the required properties of a constant-density core or an inner cusp.  Finally, in \\S~\\ref{sec:summary}, we summarize our conclusions and discuss future observations that could test more definitively whether PMN~J1632--0033 is a three-image quasar. ",
        "conclusions": "\\label{sec:summary}  For the purpose of finding central odd images, the best systems are radio-loud two-image quasars with a large magnification ratio between the two bright images.  Radio loudness is important because the radio components can be observed through the lens galaxy, whereas optical components are likely to be hidden by the lens galaxy.  Free-free absorption is potentially significant but can be overcome by observing at high frequencies.  Two-image systems are better than four-image systems, as emphasized recently by \\citet{keeton02}.  This is because the projected source position in four-image systems is usually closer to the lens center, where the central image magnification is smallest. In addition, four-image systems are subject to a larger magnification bias than two-image systems \\citep{wn93,king96,kks97,rt01,finch02}, leading to a smaller intrinsic source flux and a fainter odd image. Among two-image systems, the asymmetric doubles are favorable because the source position is close to the radial caustic, producing the brightest possible central image for a given mass distribution.  The subject of this paper, PMN~J1632--0033, is exactly this type of system: an asymmetric two-image radio lens.  For this reason, it seemed reasonable that the third component found by \\citet{winn02} could be an example of the elusive and long-sought central odd images. However, we have shown that the third component has a spectral index that differs by $3\\sigma$ from those of the two bright quasar images.  We have argued that the two most plausible interpretations are that the central component is emission from the lens galaxy, or else that it is a third image with the extra complication that its spectral index has been modified by free-free absorption.  The data are quantitatively consistent with either scenario.  Lens galaxies are rarely radio-loud, but the flux density of the component is consistent with typical radio powers of active galactic nuclei.  No central image has ever been securely identified for a radio lens, nor has free-free absorption by a lens galaxy ever been reported (with one possible exception, {\\sc class}~B0128+437; I.\\ Browne 2002, private communication), but one might expect the two phenomena to be related due to the hot and dense conditions near galaxy centers.  Given the new VLBI data, we have investigated models for the mass distribution of the lens galaxy under various assumptions.  Whether or not component C is a third image, we showed that if the mass distribution obeys a power law $\\rho\\propto r^{-\\beta}$, then it must be nearly isothermal.  If C is the lens galaxy then $1.95<\\beta<2.28$. If C is a third image then the profile is tightly constrained to be only slightly shallower: $1.89<\\beta<1.93$.  We also derived limits on a constant-density core radius, or on the break radius and inner exponent of a central cusp.  These results can be added to the growing collection of evidence that early-type galaxies have nearly isothermal mass profiles.  In the few other gravitational lens systems where it is possible to measure the radial density profile, the results favor isothermal profiles \\citep{kochanek95,cohn01,mkk01, rusin02,kt02}, although there are some possible counter-examples \\citep{tk02,ckh95}.  Moreover, the findings are consistent with studies of early-type galaxies based on stellar dynamics \\citep[see, e.g.,][]{gerhard01,rix97} and X-ray halos \\citep[see, e.g.,][]{fabbiano89}.  To further elucidate the nature of the central component and the central density profile of the lens galaxy, it would help to have a better optical image of the lens galaxy. The current HST image provides only a weak detection of the lens galaxy.  A better image could be used to test more precisely whether radio component C is coincident with the center of the lens galaxy, and to study its surface brightness profile in detail.  For example, with a better measurement of the galaxy light profile, one might be able to explore the balance between the luminous and dark matter in this system.  Even more helpful would be sensitive observations at high radio frequencies ($\\gsim$30~GHz).  As shown in Fig.~\\ref{fig:ratios}, the power-law and the free-free absorption fits diverge significantly at high frequencies, where free-free absorption becomes negligible.  For free-free absorption, the ratio \\covera~must converge to a constant value.  For lens galaxy emission, the spectrum need not follow a power law at the high frequencies, but it would require an unlikely coincidence for the intrinsic radio spectrum of the lens galaxy to turn over at high frequencies and become identical to the radio spectra of A and B.  Observations at low radio frequencies ($\\lsim$1~GHz) could also distinguish the hypotheses, in principle, but it would be harder to achieve the necessary angular resolution and signal-to-noise ratio.  High frequency observations present their own challenges, including high noise levels and susceptibility to poor atmospheric conditions, but the reward would be a definitive discovery of one of the long-sought but never-seen central images of a gravitational lens."
    },
    "0212/astro-ph0212279_arXiv.txt": {
        "abstract": "We have measured the absolute proper motions of four low-latitude, inner-Galaxy globular clusters. These clusters are: NGC 6266 (M62), NGC 6304, NGC 6316 and NGC 6723. The proper motions are on the $Hipparcos$ system, as no background extragalactic objects are found in these high-extinction regions. The proper-motion uncertainties range between 0.3 and 0.6 mas yr$^{-1}$.  We discuss the kinematics of these clusters and of three additional bulge clusters --- NGC 6522, NGC 6528 and NFC 6553 --- whose  proper motions with respect to bulge stars had been determined previously. We find that all of the clusters have velocities that confine them to the bulge region. Of the three metal poor clusters ([Fe/H] $< -1.0$), NGC 6522, and NGC 6723 have kinematics consistent with halo membership. The third cluster, NGC 6266 however, appears to belong to a rotationally-supported system. Of the four metal rich clusters ([Fe/H] $\\ge -1.0$), NGC 6304 and NGC 6553 also have kinematics consistent with membership to a rotationally-supported system. NGC 6528 has kinematics, metallicity and mass that argue in favor of a genuine Milky-Way bar cluster. NGC 6316's kinematics indicate membership to a hotter system than the bar. ",
        "introduction": "Our Galaxy's inner regions are a challenge for Galactic structure studies, as they consist of a superposition of stellar populations of various spatial, physical and kinematical properties, all of which peak in density in these regions. In order to understand the stellar populations of the inner Galaxy, one desires a tracer population that has, at the least, reliable measurements of distances, radial velocities, proper motions and metallicities, and, ideally, detailed chemical abundances and ages. Globular clusters fulfill these requirements well. Their only drawback is their small number, when compared to that of field stars.  There is still the question of whether the clusters are representative of the Galaxy's field stellar population. Many recent studies point to the fact that the globular-cluster population does resemble in many ways the old stellar field component. Based on radial velocities and metallicities it has been shown that globulars can be divided into thick disk and halo components (Zinn 1985, Armandroff 1989). Later, Minniti (1995) studied the density distribution of  metal rich clusters in the bulge region ($r_{GC} \\le 3$ kpc)  \\footnote[1]{Here, $r_{GC}$ represents the Galactocentric radius in a spherical coordinate system, and $R_{GC}$ the Galactocentric distance in a cylindrical coordinate system.}, and he concluded that most of these clusters belong to the bulge system rather than the disk. More recent distance, radial-velocity, and metallicity measurements of clusters in the bulge region together with more detailed analyses (Barbuy, Bica, \\& Ortolani 1998, C\\^{o}t\\'{e} 1999, Heitsch \\& Richtler 1999; hereafter HR99, Burkert \\& Smith 1997; hereafter BS97) suggest that some clusters may form a bar system that resembles the Milky Way bar (see Gerhard 2002 for a review of observations and characteristics of the Galactic bar).  Regarding clusters outside the bulge region but within the solar circle, Zinn (1996) shows that intermediate metal-poor clusters ($-1.8 < $ [Fe/H] $ < -0.8$) have significant rotation ($V_{rot} \\sim 60$ km s$^{-1}$) which is however, lower than that of the thick disk. This analysis included radial-velocity data alone and the typical Frenk \\& White (1980) solution, that can determine a mean rotation of a pre-selected group of clusters.  Adding tangential velocities as a piece of crucial information, Cudworth \\& Hanson (1993) and later, Dinescu, Girard \\& van Altena (1999; Paper III) have unambiguously shown that there is a metal-weak thick disk component in the cluster population, much like that seen in the local field stellar population (see Beers {\\it et al.} 2002 for the most recent study). And, the inner-halo rotation detected by Zinn (1996), was also found in the studies that included tangential velocities (Paper III). This rotation is qualitatively consistent with the flattened shape of the inner halo as determined from field stars (e.g., Hartwick 1987, Chiba \\& Beers 2001, Siegel {\\it et al.} 2002). To summarize, the globular-cluster system of our Galaxy appears to resemble the old stellar field component.  We set out to measure tangential velocities (i. e., absolute proper motions) for a sample of fifteen clusters (Dinescu, Girard, \\& van Altena 2002) primarily located within the solar circle, in order to more accurately assign clusters to the components of the inner Galaxy, and to better understand these components as regions where old and massive stellar systems were produced/captured.  In this paper we present the first results for four globular clusters located in the bulge region: NGC 6266 (M 62), NGC 6304, NGC 6316, and NGC 6723. We discuss the kinematics of our four globular clusters and of three additional clusters (NGC 6522, NGC 6528, and NGC 6553) located in the bulge region that have previous proper-motion measurements, in light of the Galactic components seen in the bulge region.  The proper-motion measurements are described in Section 2, the velocity results are presented in Section 3, and discussed in Section 4, and a brief summary is presented in Section 5. ",
        "conclusions": "All of the clusters under discussion reside within $\\sim 3$ kpc of the GC (Tab. 4). This is a very complex region dominated by a bar/bulge where a traditional kinematical distinction between the disk and the bar/bulge may not be as meaningful as in more distant regions from the GC, where  disk and halo populations for instance are easily separated. Dynamical criteria allow the existence of a kinematical disk only outside the corotation radius (see e. g.,  Pfenniger \\& Friedli 1991). N-body simulations show that bars form easily as a gravitational instability in a cold disk, and, in order to survive, they have to rotate (Pfenniger 1999). Bars then, can be regarded as ``thickened'' elliptical disks. They preserve the ``memory'' of the original rotating axisymmetric disk, and therefore one may be able to kinematically  distinguish between the bar/bulge and a pressure-supported system such as the halo. Presumably, a halo particle will remain a halo particle, while a disk particle will become a bar/bulge particle inside the corotation radius. In N-body simulations, there is evidence of some resonant halo stars trapped in the bar, in particular at corotation (Athanassoula 2002); however, for our study that considers very small numbers, we believe that it is unlikely to find a halo globular cluster on a resonant orbit with the Galactic bar.  A separation between bar and bulge is more ambiguous given the variety of processes that may contribute to bulge formation such as bar growth, bar dissolution, accretion and/or merging with satellites.  In our Galaxy, evidence for a bar comes from observations of motions of atomic and molecular gas, from near-infrared photometry (NIR), IRAS sources and clump stars counts, and stellar kinematics (see e.g., Gerhard 1996 for a review). The characteristics of the Milky Way bar are summarized in Gerhard (2002). The bar has a semimajor axis a = $3.1 - 3.5$ kpc, an axis ratio of 10:3-4:3, and an orientation angle (angle between the semimajor axis and the Sun-GC line) of $\\phi_{bar} = 15\\arcdeg - 35\\arcdeg $. The pattern speed is $\\Omega_b = 50 - 60$ km s$^{-1}$ kpc$^{-1}$, and the corresponding corotation radius is $R_{CR} = 3.3 - 5 $ kpc. The Galactic bar seems to be a fairly massive system (M$_{bar} \\sim 10^{10}$ M$_{\\odot}$; Stanek et al. 1997, Weiner \\& Sellwood 1999), of an age of up to $\\sim 6 - 7$ Gyr (as determined from Carbon stars, Cole \\& Weinberg 2002 and references therein). Given these considerations, can we speak of a globular-cluster bar system? Were these clusters trapped into bar-like orbits from an initially disk-like configuration, as dissipationless N-body simulations show? Or, were these clusters produced within the bar, if the bar formed from gas rather than from a stellar distribution? BS97 argued that low-mass (i.e. integrated absolute magnitude M$_V \\ge -7.9$), metal-rich ([Fe/H] $ > -0.8$) clusters show spatial distribution and kinematics --- as derived from radial velocities --- consistent with those of the Milky Way bar (see also C\\^{o}t\\'{e} 1999 and references therein). Given the large mass of the bar, and its upper age limit, it is quite conceivable that low-mass, metal-rich clusters formed within the bar. Dating the clusters suggested by BS97 as belonging to a bar-like configuration could be very instructive.  In light of the above, we proceed to discuss the results for the seven clusters under study, without attempting a detailed orbit integration that we defer to a future paper.   In Table 5 we list the cluster metallicity [Fe/H], integrated absolute magnitude M$_V$ (H96), total velocity, escape velocity from the bulge, and orbit inclination. The orbit inclination is $\\Psi = 90\\arcdeg -$ sin$^{-1}$(L$_z$/L), where L$_z$ is the angular momentum perpendicular to the Galactic disk, and L is the total angular momentum, at the present time. In contrast to Paper III, this is not an orbit-averaged quantity, but an instantaneous one. The bulge potential is the Hernquist spheroid described in Johnston, Spergel \\& Hernquist (1995; JSH95). This potential is a simple description of the inner Galaxy, and it is used only to estimate the escape velocity and not to derive orbits. As total velocities are smaller than escape velocities, all of the clusters are confined to the bulge region.   In Figure 4 we show the locations and motions of the clusters in the Galactic plane (top panel) and perpendicular to the Galactic plane (bottom panel). The open symbols represent the metal poor clusters, the filled symbols the metal rich ones (see following subsections). The dashed line represents the bulge region; it has a radius of 2.7 kpc, which corresponds to an angle of $20\\arcdeg$ on the sky from the GC (from COBE and IRAS maps, and $R_{\\odot} = 8.0$ kpc). The ellipse represents the bar. Its characteristics are those determined by Bissantz  \\& Gerhard (2002) from a model that is constrained by the dereddened COBE/DIRBE L-band luminosity distribution and by the distribution of red clump stars in some bulge fields. The bar has a semimajor axis of 3.5 kpc, an axis ratio of 10:3:3, and an orientation angle of $20\\arcdeg$.   \\subsection{Metal Poor ([Fe/H] $< -1.0$) Clusters}   NGC 6522 and NGC 6723 have locations and velocity components that produce highly inclined orbits (Tab. 5). Their kinematics suggest these clusters are members of the halo system. Their low metallicities are in agreement with this assignment (e. g., Armandroff 1989). We note that, within the uncertainties of their distance estimates, NGC 6723 and NGC 6522 might be positioned on either the far side or near side of the GC, in which case the velocity components $\\Pi$ and $\\Theta$ change signs. Thus retrograde orbits become prograde and vice versa; however this does not alter the large inclination of their orbits. In other words, the sense of rotation with respect to the disk is of little significance when orbits are nearly polar (Tab. 5).  NGC 6266 has a large rotation velocity, while the other two components are relatively low (Tab. 4). This velocity vector, combined with a galactocentric distance $R_{GC} = 1.4$ kpc (Tab. 4) indicate that NGC 6266 belongs to a rotationally supported system. Its rotation velocity is somewhat larger than that derived for a solid body rotating with the Galactic bar's pattern speed (Tab. 4). As the notion of a kinematical stellar disk is inappropriate at these galactocentric distances, we can think of NGC 6266 as a cluster that belonged to a disk system before the bar formed, rather than to a pressure-supported system such as the halo. In this sense, NGC 6266 reminds us of thick-disk, metal-poor clusters such as NGC 6626, NGC 6752, and perhaps NGC 6254 (see Paper III). These clusters represent the metal-weak thick disk (MWTD), found to be very prominent in the field stellar component within 1 kpc of the Sun (Beers {\\it et al.} 2002 find a fraction of 30-40\\% MWTD at metallicities $<$ -1.0). The novelty provided by the clusters is that this component can be probed at smaller galactocentric radii than local samples.  \\subsection{Metal Rich ([Fe/H] $ > -1.0$) Clusters}  Based on spatial distribution, radial velocities and masses, BS97 have proposed a new classification of clusters with [Fe/H] $ > -0.8$. The high-mass system  (M$_V < -7.9$) is thought to be representative of the bulge. It has an angular velocity $\\omega = 15$ km s$^{-1} $ kpc$^{-1}$, and a line-of-sight velocity dispersion  $\\sigma_{los} = 56$ km s$^{-1}$, assuming  a solid-body rotation in a Frenk \\& White (1980) type solution. The low-mass system (M$_V \\ge -7.9$) is a sum of two components: a bar-like component, and a thick-disk-like component. The bar-like  component has an angular velocity ($\\sim 59$ km s$^{-1} $ kpc$^{-1}$) similar to traditional values of Milky Way bar's pattern speed, and a  $\\sigma_{los} = 55$ km s$^{-1}$. The thick-disk-like component resides at $R_{GC} \\sim 5$ kpc, and it has typical thick-disk kinematics (see Table 1 in BS97).   According to BS97, by mass and metallicity, NGC 6316 is a bulge cluster (Table 5). If we use the angular velocity of the bulge globular-cluster system defined in BS97 to predict the rotation velocity at NGC 6316's location, we obtain that NGC 6316's retrograde velocity is at 2.2$\\sigma$ of that predicted by the solid body. To estimate $\\sigma$ we have used the uncertainty in the $\\Theta$ component (Tab. 4) and we have taken the velocity dispersion along $\\Theta$ to be equal to the line-of-sight velocity dispersion derived from the solid-body solution for all of the candidate bar clusters in BS97. The constant-rotation velocity solution in BS97  places NGC 6316 within 1.6$\\sigma$ of the mean of the globular-cluster bulge system defined by BS97, as the line-of-sight velocity dispersion is $\\sim 70$  km s$^{-1}$. A similar comparison with the BS97 bar system, places NGC 6316 at 4.3$\\sigma$ from the rotation velocity expected from the solid body. Thus, according to its velocity, NGC 6316 can be thought of as a bulge cluster rather than a bar cluster, and this is in agreement with its classification according to metallicity and mass.   Both NGC 6304 and NGC 6553 have large rotation velocities and low to moderate radial and vertical velocity components (Tab. 4). The $\\Theta$ components are somewhat larger than the rotation velocities expected from a solid body that has an angular velocity equal to the Galactic bar's pattern speed (Tab. 4). The large $\\Theta/V'$-velocity component of NGC 6553 may be due to the assumption that the reference stars (bulge stars) in the proper-motion solution have zero velocity (see Section 3.1). The Galactic longitude of this cluster is larger than that of NGC 6522, and 6528, therefore this assumption may not be correct. In fact, at a longitude of $ \\sim 5\\arcdeg$ (Table 3), the bulge stars are expected to have a rotation of 45 km s$^{-1}$ (Tiede \\& Terndrup 1999 measure a projected mean rotation of 9 km s$^{-1}$ deg$^{-1}$). Thus, the rotation-velocity component should, more likely, be $V' \\sim 180$ km s$^{-1}$, and correspondingly, $\\Theta = 175$  km s$^{-1}$, a value that is much closer to NGC 6304's rotation velocity. Masses and metallicities assign these clusters to the bar system of BS97 (NGC 6553 is slightly more massive than the BS97 limit; see Tab. 5). Their kinematics, at face value, are not inconsistent with bar membership, especially if we must abandon the idea of disk-like orbits inside the corotation radius. These clusters could very well be members of the former disk/thick-disk transformed into a bar in the inner regions. They can be thought of as the extension of the 5-kpc thick disk clusters in BS97 into the inner Galactic regions. What is intriguing is that NGC 6266, a massive, metal-poor cluster, is kinematically very similar to NGC 6304 and NGC 6553 (Tab. 4, 5; Fig. 4). The slight difference is that NGC 6266, while currently located farther from the Galactic plane than NGC 6304 and NGC 6553 (Fig. 4), moves away from the plane with the largest velocity of the three clusters (Tab. 4). Strictly speaking, this implies that NGC 6304 and 6553 are in a more flattened system than NGC 6266.  NGC 6528 qualifies as a bar cluster according to BS97 (Tab. 5). Its $\\Theta$ velocity component is in very good agreement with the rotation of a solid body that has the Galactic bar's pattern speed (Tab. 4). The very low vertical velocity together with the close distance to the Galactic plane (Fig. 4) indicate that the cluster will spend its time near the Galactic plane, while the large $U$-velocity component, indicates a radial orbit. We note that NGC 6528's distance has a wide range of values, even from more recent determinations (e.g., Feltzing \\& Johnson 2002 determine a distance of 7.2 kpc, while HR99 give a distance of 9.3-10.7 kpc from isochrone fitting, and from the horizontal-branch metallicity relation, respectively). Even if, according to distance errors, NGC 6528 is displaced from the far side of the bulge to the near side, the $U$-velocity component will remain practically the same (while the $\\Pi$ component changes sign, but $\\Theta$ does not,  see Fig. 4). This will leave the character of the orbit unchanged.  These kinematical arguments together with the fact that NGC 6528 is, arguably, the most metal rich Galactic globular cluster (Feltzing \\& Johnson 2002), make a strong case for a globular formed within the bar.   We note that, the two clusters that are located outside the bar limits, NGC 6316 and 6723 (Fig. 4), also have kinematics inconsistent with that of the bar. Among the clusters located within the bar limits, only NGC 6522 --- a metal poor cluster --- seems to have kinematics inconsistent with that of the bar."
    },
    "0212/astro-ph0212086_arXiv.txt": {
        "abstract": "Many of the properties of anomalous X-ray pulsars (AXPs) and soft gamma-ray repeaters (SGRs) are still a matter of much debate, as is the connection (if any) between these two groups of sources.  In cases where we can identify the supernova remnant (SNR) associated with an AXP or SGR, the extra information thus obtained can provide important constraints as to the nature of these exotic objects.  In this paper, I explain the criteria by which an association between a SNR and an AXP/SGR should be judged, review the set of associations which result, and discuss the implications provided by these associations for the ages, environments and evolutionary pathways of the AXPs and SGRs. There are several convincing associations between AXPs and SNRs, demonstrating that AXPs are young neutron stars with moderate space velocities. The lack of associations between SGRs and SNRs implies that the SGRs either represent an older population of neutron stars than do the AXPs, or originate from more massive progenitors. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/astro-ph0212014_arXiv.txt": {
        "abstract": "Applications of the forthcoming \\GAIA\\ satellite to the study of variable phenomena (microlensing and supernovae) are discussed. ",
        "introduction": "\\GAIA\\ is an astrometric satellite which has recently been approved by the European Space Agency for launch in about 2010. It will measure the angles between objects in fields that are separated on the sky by about a radian. Data will stream continuously at 1~Mbps from \\GAIA's three telescopes, providing information on the positions of the billion or more astrophysical objects brighter than 20th magnitude. The motion of objects across the sky, due to both their space motion and their parallactic motion and the variability in the brightnesses of objects in 15 wavebands will be measured because each object is observed at least 150 times during the 5 year mission lifetime. From the raw time series, a three-dimensional movie of the motions of stars in the Galaxy will be synthesized. Just as the human brain interprets signals from the retina in terms of moving three-dimensional objects, so \\GAIA\\ will interpret signals from its sensors in terms of moving astrophysical objects. \\GAIA's job is dramatically the more difficult because (1) \\GAIA's sensors return time-series rather than images, (2) the global adjustment of the myriad of instantaneous positions is a computational task of huge complexity, (3) \\GAIA\\ has to correct for gravity-induced distortions of space-time, and (4) \\GAIA\\ sees in 15 colours rather than three.  This article presents highlights from the five-year \\GAIA\\ movie show. The alert despatcher in the \\GAIA\\ mission will provide forewarning of many kinds of bursting and variable phenomena. We discuss two applications in detail here. First, the astrometric microlensing signal will allow us to take a complete inventory of all objects -- no matter how dark -- in the solar neighbourhood.  Second, the catalogue of supernova (SN) detections will be the largest ever taken. It will provide opportunities both for follow-ups in other wavebands and numerous examples of scarce phenomena (like SNe II-L or SNe Ib/c). ",
        "conclusions": "The \\GAIA\\ satellite is scheduled for launch in about 2010. It will determine the positions, velocities and astrophysical nature of over a billion stars distributed throughout the Milky Way Galaxy and into the Local Group. Each object will be monitored between 150 and 400 times during the 5 year mission (depending on ecliptic latitude). The resulting catalogue will be one of the largest astrophysical datasets ever taken. The positions and velocities of stars, together with the changing brightness of variable sources (e.g., bursting and eruptive stars, supernovae, killer asteroids, eclipsing variables, microlensed stars), will all be synthesized into a three-dimensional movie show.  This article has discussed two applications of real-time detection of variable sources with the \\GAIA\\ satellite. \\GAIA\\ will provide the best determination of the local mass function of stars down to brown dwarf mass scales -- which is important for reckoning the baryonic and particle dark matter content of the solar neighbourhood. This can be inferred from the astrometric microlensing signal of nearby stars. \\GAIA\\ will provide the largest sample of nearby supernovae ever acquired -- which is crucial for assessing how the scatter in properties affects their use as cosmological distance indicators. This will provide opportunities both for follow-ups in other wavebands (e.g., X ray and gamma-ray) and with other detectors (e.g., neutrino, gravitational waves).  Further applications are given in references~\\cite{5} and~\\cite{14}."
    },
    "0212/astro-ph0212222_arXiv.txt": {
        "abstract": "We study galaxy pairs in the field selected from the 100 K public release of the 2dF galaxy redshift survey. Our analysis provides a  well defined sample of 1258 galaxy pairs, a large database suitable for statistical studies of galaxy interactions  in the local universe, $z \\le 0.1$. Galaxy pairs where selected by radial velocity ($\\Delta V$) and projected separation ($r_{\\rm p}$) criteria determined by analyzing the star formation activity within  neighbours. We have excluded pairs in high density regions by removing galaxies in groups and clusters. We analyze the star formation activity in the pairs as a function of both  relative projected distance and relative radial velocity. We found power-law relations for the mean star formation birth parameter and equivalent widths of the galaxies in pairs as a function of $r_p$ ad $\\Delta V$. We find that star formation in galaxy pairs is significantly enhanced over that of isolated galaxies with similar redshifts in the field for  $r_p < 25$ kpc and $\\Delta V < 100$ km/s. We detected that when compared to isolated galaxies of similar luminosity and redshift distribution, the effects of having a companion are more significant on the star formation activity of bright galaxies in pairs, unless the pairs are formed by similar luminosity galaxies. In this case, the star formation is enhanced in both components. The ratio between the fractions of star forming galaxies in pairs and in  isolation is a useful tools to unveil the effects of having a close companion. We found that about fifty percent of galaxy pairs do not show signs of important star formation activity (independently of their luminosities) supporting the hypothesis that the internal properties of the galaxies play a crucial role in the triggering of star formation by interactions. ",
        "introduction": "The  actual cosmology paradigm for galaxy formation assumes that structure forms by hierarchical aggregation. In such scenarios galaxy interactions are  frequent and have a crucial role in determining galaxy properties. The formation of galaxies in these schemes can be successfully analyzed by using numerical simulations which show that, as the systems are assembled, mergers and interactions can trigger star formation with efficiencies that seem to depend mainly on the internal structure of the systems. By using pre-prepared mergers, Barnes \\& Hernquist (1996) showed how the gas component experiences torques originated in the companion, increasing its gas density and triggering a starburst during the orbital decay phase of the satellite. These starbursts are fed by gas inflows  tidally induced if the axisymmetrical character of the potential well is lost during the interaction. The stability of the systems can be assured by a dominating central mass concentration.  A second starburst could be generated at the actual fusion of the baryonic cores. Cosmological hydrodynamical simulations showed that  these processes take place in the formation of galactic systems in a hierarchical aggregation scenario  in a similar way to that shown by pre-prepared mergers. Tissera (2000) found a correlation between mergers and starbursts. Recently, Tissera et al. (2002, Tis02) showed that the effects of interactions are different at different stages of evolution of the systems being  more efficient at higher $z$ when the systems are in early stages of evolution, and consequently, their potential well could be shallower. If the structure in the Universe formed in consistency with a hierarchical scenario, then, the proximity of  a companion could affect the mass distributions and  trigger gas inflows producing an enhancement of the star formation (SF) activity.  In the local Universe, observations show that mergers and interactions of galaxies affect star formation in galaxies (e.g., Larson \\& Tinsley 1978; Donzelli \\& Pastoriza 1997; Barton, Geller \\& Kenyon 1998). It has also been shown that the merger rates (e.g., Woods, Fahlman, Richer 1995; Le Fevre et al. 2001; Patton et al. 2002) and the star formation activity of galaxies increase with redshift, suggesting a change in the impact of interactions on the SF process as galaxies evolve. Although the relevance of these violent events on the formation of the structure and their SF history is now widely accepted, many questions remain to be answered. For example, as it has been reported by other authors (e.g., Petrosian et al. 2002), many interacting systems show weak star formation activity suggesting that the particular internal conditions within these systems  may be needed to trigger star formation.  An insight on the nature of interactions can be obtained from studies of close pairs of galaxies in projection. Physical pairs must have similar redshifts, where relative velocities affect distance interpretation and therefore the true relative separations. Yee \\& Ellingson (1995) and Patton et al. (1996) adopt a minimum projected separation of $20 h^{-1} kpc$ to define close pairs and they find no significant differences between mean properties of paired and isolated galaxies although those which appear to be undergoing interactions or mergers have strong emission line and blue rest-frame colours. A similar result by Zeff \\& Koo (1989) indicate that in some systems colours correspond to recently enhanced star formation with an overall  distribution consistent with  field galaxies. Kennicutt et al. (1987) examined the $H_\\alpha$ equivalent width, UBV colours and far-infrared flux of a complete sample of local pairs of galaxies and found that these close pairs exhibit a general trend of enhanced star formation and nuclear activity but with a wide dispersion about the mean behavior. Barton et al. (1998) analyzed a sample of approximately 250 pairs of  galaxies showing a correlation between their radial and velocity separations and the $H_{\\alpha}$ equivalent width.  Although these authors found a clear correlation, their pair sample is still small to carry out a thoughtful statistical study of the effects of interactions and their cosmological evolution.  The 100K public data release (Colles et al. 2001) of approximately 100000 galaxies of the 2dF Galaxy Redshift Survey (2dFRS) opens the possibility to study galaxy interactions by selecting the largest galaxy pair catalog hitherto constructed and analyzing their properties. The 2dF public data provide the redshift, magnitude and spectral types ($\\eta$) of galaxies with $z \\le 0.15$ (Madgwick et al. 2002a, hereafter M02). Madgwick et al. (2002b) found a tight correlation between $\\eta$ and the birthrate parameter $b$.  In this work, we present a catalog of galaxy pairs in the field constructed from the  2dF public release data and analyze their properties in comparison to isolated galaxies in the field. For this purpose we will exclude those galaxies that belong to groups as defined by Merch\\'an \\& Zandivarez (2001, hereafter MZ01). Pairs in high density environments will be analyzed in a forthcoming paper. Here, we aim at answering questions such as which type of galaxies are preferentially found in pairs, how the star formation varies among them according to their luminosities, etc. We also intend to statistically determine a critical  relative radial velocity  and spatial separation for the classification of galaxy pairs. As it has been  pointed out in previous works, galaxy pairs can be classified as interacting or close pairs. Interacting galaxy pairs are those that show explicit signs of interactions such as tidal tails. Close pairs are defined according to velocity difference and spatial separation criteria (e.g., Barton et al.2002; Patton et al. 2002). In this work, galaxy pairs will be statistically selected by applying both a velocity and separation criteria which will be determined according to the star formation  activity of neighbouring galaxies.  Because the public release 2dF catalog has different selection effect problems, control samples sharing the same selection effects will be constructed,  focusing the analysis on the statistical differences between galaxies in pairs and isolated ones in the 2dFRS as a useful procedure to unveil  the effects of galaxy interactions on the star formation process.  In Section 2 we will defined the pair selection criteria , field galaxy pair catalog and the comparison samples. Section 3 is focused on the analyses of the star formation properties of galaxies in pairs and  Section 4 summarizes the main conclusions. ",
        "conclusions": "We carried out a statistical analysis of 1258 galaxy pairs in the field with $z \\le 0.1$  selected from the  2dF public release data. Our study is centred on the star formation enhancement of the pair members with respect to isolated galaxies in the field with the same redshift distribution and luminosities by analysis the dependence of the mean birth rate paremeters and the fractions of star forming galaxies. Particularly, the latter results to be an important tool to unveil the effects of having a companion since it is less affected by interpolers, non SF pairs, etc.   We can summarize our results in the following main conclusions.  1. We detect a significant correlation between the star birth rate parameter $<b>$ and both, projected spatial separation $r_p$ and  relative radial velocity $\\Delta V$ for galaxy pairs. For $r_p < 25 $kpc and $\\Delta V < 100$ km/s we obtain a substantial star formation enhancement with respect to the isolated control sample. The $\\Delta V$ dependence is less pronounced although we find a systematic increase of star formation activity for decreasing relative velocity.    2.  A significant percentage of galaxies ($\\approx 45 \\%$ ) which satisfy the pair selection criteria but that have negligible star formation has been found. This result suggests that it is not only proximity that may be playing a role in triggering star formation but the internal structure of the galaxies could also  be a crucial factor as it has been found in numerical simulations.  3.  There is no an overall dependence of the mean star formation enhancement $<\\beta> $ on luminosity for galaxy pairs at $z \\leq 0.10$. A nearly constant value  $<\\beta> \\simeq 1.2  $  for the total FGP catalog and $<\\beta> \\simeq 2 $  for the close GP subsample have been measured. However, the fractions of galaxies in pairs that have a higher SF activity than the  corresponding averaged one  of isolated galaxies, show a clear excess for the bright members.   4. We divide  the close pair sample in  minor ($L_2/L_1 < 0.5$) and major ($L_2/ L_1 > 0.5$) merger candidates according to their relative luminosities. We found that bright components in minor merger candidates show higher probability to have enhanced SF by tidal interactions than isolated galaxies and from larger projected distances than the faint components. In the case of major mergers both components show comparable star formation enhancements and with a similar projected distance dependence."
    },
    "0212/astro-ph0212478_arXiv.txt": {
        "abstract": "A two-week campaign of high-resolution imaging of the centre of M15 using the tip/tilt correcting TRIFFID-2 camera, followed by analysis with ISIS (image matching \\& subtraction software), has produced the most sensitive survey to date of variable stars in the central $\\sim$1$\\arcmin$$\\times$1$\\arcmin$ - a total of 48 were detected -, and constrains the size of the dwarf nova population. ",
        "introduction": "M15 is a massive core-collapsed cluster, with an enormous stellar density in the central $\\sim$1$\\arcmin$$\\times$1$\\arcmin$, reflected in the scarcity of photometric variable star detections in this field. However, Ferraro \\& Paresce (1993) used the HST/FOC to identify 19 variable candidates. Our earlier work (Butler et al. 1998), using a TRIFFID/MAMA camera similar to that described here, resulted in light curves and periods for all but four of these, with a further 16 suspected new variables.  We made new observations in \\emph{B} and \\emph{V} simultaneously with the TRIFFID-2/MAMA camera on the 1m JKT (La Palma), over 12 nights in July 1997, with seeing of 0$\\farcs$7-0$\\farcs$8. The 2-d photon-counted data were first sharpened post-exposure by tip/tilt correction at 1ms time resolution.  The ISIS (Alard \\& Lupton 1998) image matching and difference imaging package was used to search for frame-to-frame variability.  All star-like objects 6-sigma above the local background of the the {\\it median} of the set of difference images were selected as candidate variables (see Figure 1). PSF-fitting photometry was performed at accurate positions for the candidate variables in the difference images. The resulting relative fluxes were converted to absolute fluxes by performing profile-fitting photometery on the reference image. The instrumental magnitudes were zeropointed against photometry of HST/WFPC2 archival images in the F555W and F439W filter system. Accurate periods (and hence lightcurves) were obtained for the set of magnitude measurements for all candidate variables using the Phase Dispersions Minimisation technique (PDM, Stellingwerf 1978). Simulations performed using an artificial image-set modeled on our M15 observations indicate that we are sensitive to RR Lyrae variables with an amplitude swing of 0.15 mag (at 80\\% detection probability). We also increased our sensitivity to dwarf novae (DNe) by summing all images per night prior to ISIS analysis.    \\begin{figure} \\plottwo{otuairisg_fig1.ps}{otuairisg_fig2.eps} \\caption{Left: one V-band image of M15. Right: median-difference image, clearly showing the variable candidates (numbered).} \\end{figure} ",
        "conclusions": ""
    },
    "0212/astro-ph0212152_arXiv.txt": {
        "abstract": "{ Fast stellar winds can sweep up ambient media and form bubbles. The evolution of a bubble is largely controlled by the content and physical conditions of the shocked fast wind in its interior. This hot gas was not clearly observed until the recent advent of {\\it Chandra} and {\\it XMM-Newton} X-ray observatories. To date, diffuse X-ray emission has been unambiguously detected from two circumstellar bubbles blown by WR stars, four planetary nebulae, and two superbubbles blown by young clusters. Model fits to the X-ray spectra show that the circumstellar bubbles are dominated by hot gas with low temperatures ($\\le 3\\times 10^6$ K), while the interstellar bubbles contain significant fractions of hotter gas ($\\ge 5\\times10^6$ K). In all cases, large discrepancies in the X-ray luminosity are found between observations and conventional models of bubbles. Future theoretical models of bubbles need to re-examine the validity of heat conduction and take into account realistic microscopic processes such as mass loading from dense clumps/knots and turbulent mixing.  {\\it Chandra} ACIS-S observation of NGC\\,6888 will shed light on these astrophysical processes.}    \\addkeyword{H\\,II regions} \\addkeyword{ISM: bubbles} \\addkeyword{H~II regions} \\addkeyword{Planetary nebulae} \\addkeyword{Stars: Mass loss} \\addkeyword{Stars: Wolf-Rayet stars} \\addkeyword{X-rays}  \\begin{document} ",
        "introduction": "In 1965, Johnson \\& Hogg reported three shell nebulae, NGC\\,2359, NGC\\,6888, and the nebula around HD\\,50896, and suggested that these nebulae were formed by interactions between the mass ejected by the central Wolf-Rayet (WR) stars and the ambient interstellar gas. In modern terminology, these shell nebulae are called ``bubbles\", the mass ejected by a WR star is ``fast stellar wind\", the shell nebula around HD\\,50896 has been cataloged as ``S\\,308\", and the ambient medium is designated ``circumstellar medium\", i.e., ejected stellar material.  The earliest theoretical treatments of interactions between fast stellar winds and the interstellar medium (ISM) were motivated by the central cavities and high-velocity motions observed in \\ion{H}{2} regions (Mathews 1966; Pikel'ner 1968; Pikel'ner \\& Schcheglov 1969; Dyson \\& de Vries 1972). Models of wind-ISM interaction were constructed specifically for shell nebulae around WR stars by Avedisova (1972), and a numerical calculation of wind-ISM interaction was provided by Falle (1975).  In 1974, Jenkins \\& Meloy reported {\\it Copernicus} satellite observations of 32 early-type stars, which revealed ubiquitous shallow interstellar \\ion{O}{6} absorption.  To explain this \\ion{O}{6} absorption, Castor, McCray, \\& Weaver (1975) modeled wind-ISM interaction and coined the term ``interstellar bubble\". Weaver et al.\\ (1977) elaborated on this model and gained popularity because of their detailed description of the physical conditions and structure of a bubble and their specific characterization of X-ray emission and \\ion{O}{6} column density expected from a bubble interior.  High-quality, high-resolution X-ray and far UV observations were not possible until {\\it Chandra X-ray Observatory}, {\\it XMM-Newton X-ray Observatory}, and {\\it Far Ultraviolet Spectroscopic Explorer (FUSE)} satellites were launched. It is finally possible to observe the hot gas in bubble interiors and critically examine the bubble models. The analysis of {\\it FUSE} observations is complex, and it is pre-mature to conclude on the \\ion{O}{6} absorption from bubbles.  Therefore, in this paper we will concentrate on only X-ray observations of bubbles and compare them to model predictions.   \\begin{table*}[!t]\\centering \\newcommand{\\DS}{\\hspace{6\\tabcolsep}} % \\setlength{\\tabnotewidth}{0.9\\textwidth} \\setlength{\\tabcolsep}{1.33\\tabcolsep} \\tablecols{5} \\setlength{\\tabcolsep}{1.5\\tabcolsep} \\caption{X-ray Observing Facilities} \\begin{tabular}{clccc} \\toprule X-ray    & ~Mission  & Imaging & Angular    &  Energy \\\\ Observatory  & Duration & Spectrometer\\tabnotemark{a}  & Resolution & Range (keV)\\\\ \\midrule {\\it Einstein}  & 1978-1981  &  IPC          &  120$''$ & 0.2-3.5 \\\\ {\\it ROSAT}     & 1990-1998  &  PSPC         &  30$''$  & 0.1-2.4 \\\\ {\\it ASCA}      & 1993-2001  &  SIS          &  150$''$ & 0.4-10  \\\\ {\\it Chandra}   & 1999-pres. &  ACIS         &  1$''$   & 0.1-10  \\\\ {\\it XMM-Newton}& 1999-pres. &  EPIC         &  10$''$  & 0.1-15  \\\\ \\bottomrule \\tabnotetext{a}{IPC: Imaging Proportional Counter; PSPC: Position Sensitive Proportional Counter; SIS: Solid-State Imaging Spectrometer; ACIS: Advanced Camera for Imaging and Spectroscopy; EPIC: European Photon Imaging Camera} \\end{tabular} \\end{table*} ",
        "conclusions": ""
    },
    "0212/astro-ph0212364_arXiv.txt": {
        "abstract": "We study the $slow$ neutron capture process ($s$ process) in Asymptotic Giant Branch (AGB) stars using three different stellar evolutionary models computed for a 3 \\msb and solar metallicity star. First we investigate the formation and the efficiency of the main neutron source: the \\ct($\\alpha$,n)$^{16}$O reaction that occurs in radiative conditions. A tiny region rich in \\ctb (the \\ctb {\\it pocket}) is created by proton captures on the abundant \\cdb in the top layers of the He intershell, the zone between the H shell and the He shell. We parametrically vary the number of protons mixed from the envelope. For high local protons over \\cdb number ratio, p/\\cdb $\\gsim$ 0.3, most of the \\ctb nuclei produced are further converted by proton capture to \\nq. Besides, \\nqb nuclei represent a major neutron poison. We find that a linear relationship exists between the amount of \\cdb in the He intershell and the maximum value of the time-integrated neutron flux. Then we generate detailed $s$-process calculations on the basis of stellar evolutionary models constructed with three different codes, all of them self-consistently finding the third dredge up, although with different efficiency. One of the codes includes a mechanism at each convective boundary that simulates time-dependent hydrodynamic overshoot. This mechanism depends on a free parameter $f$, and results in partial mixing beyond convective boundaries, the most efficient third dredge up and the formation of the \\ctb pocket. For the other two codes an identical \\ctb pocket is introduced in the post-processing nucleosynthesis calculations. The models typically produce enhancements of heavy elements of about two orders of magnitude in the He intershell and of up to one order of magnitude at the stellar surface, after dilution with the convective envelope, thus generally reproducing spectroscopic observations. The results of the cases without overshoot are remarkably similar, pointing out that the important uncertainty in \\s-process predictions is the \\ctb pocket and not the intrinsic differences among different codes when no overshoot mechanism is included. The code including hydrodynamic overshoot at each convective boundary finds that the He intershell convective zone driven by the recurrent thermal instabilities of the He shell (thermal pulses) penetrates the CO core, producing a He intershell composition near to that observed in H-deficient central stars of planetary nebulae. As a result of this intershell dredge up the neutron fluxes have a higher efficiency, both during the interpulse periods and within thermal pulses. The $s$-element distribution is pushed toward the heavier $s$-process elements and large abundances of neutron-rich isotopes fed by branching points in the $s$-process path are produced. Several observational constraints are better matched by the models without overshoot. Our study need to be extended to different masses and metallicities and in the space of the free overshoot parameter $f$. ",
        "introduction": "When helium is exhausted in the center of low- to intermediate-mass stars ($M \\lsim$ 8 \\ms) the core becomes degenerate \\citep{paczynskia} and burning processes do not occur there anymore. Instead, H and He burn alternately in shells while the star ascends the Asymptotic Giant Branch (AGB). In this phase the stellar structure consists of, from the center outward, a degenerate CO core, the He-burning shell, a thin (10$^{-2} - 10^{-3}$ \\ms) zone between the H shell and the He shell (hereafter He intershell), the H-burning shell, and a large convective envelope, which suffers from strong stellar winds. When the envelope mass is reduced to less than $\\sim 10^{-3}$ \\msb the star leaves the AGB, a Planetary Nebula is formed and the CO core cools towards the White Dwarf phase. In 10-20\\% of cases the residual H-rich envelope is completely removed by processes associated with a late He-shell instability during the pre-White Dwarf phase \\citep{herwigbld}, leading to the formation of the H-deficient central stars of planetary nebulae.  During the AGB the H shell dominates the energy production for most of the time. Hydrogen is transformed into He, the He intershell consequently grows and is progressively compressed and heated until the temperature and density become high enough that He burning is triggered in the bottom layers. The thermonuclear runaway, also known as thermal instability, or {\\it thermal pulse}, \\citep{schwarzschild} generated by this sudden release of energy causes almost all the He intershell to become convective (we refer to this as the {\\it pulse driven convective zone}, PDCZ) while the envelope expands and the H shell cools. Within the PDCZ partial He burning produces large amounts of carbon. The PDCZ quenches after a time of a few ten to a few hundred years, while He burning continues radiatively for another few thousand years, during which the H shell is inactive. Then, the envelope contracts and shell H burning starts again. The cycle is repeated for a few up to possibly 100 times, the total number of thermal pulse  depending on the initial stellar mass and the mass loss rate, with thermal periods of a few 10$^{3} up to 10^{5}$ yr. Recent models of the AGB phase have been computed among others by \\citet*{hollowell,boothroyda,boothroydb,boothroydc,boothroydd,lattanzioa, lattanziob,lattanzioc,vassiliadis,blocker,forestini,stranieroc,wagenhuber,mazzitelli,mowlavi,herwig}.  As illustrated in Fig.\\ \\ref{fig1}, after a limited number of thermal pulses, when the mass of the H-exhausted core is above $\\sim 0.6$ \\ms, the convective envelope penetrates in the top layers of the He intershell bringing to the surface newly synthesized He, C and elements produced by neutron captures. This recurrent phenomenon is called {\\it third dredge up} (TDU). The TDU is responsible for the carbon enrichment shown by MS, S and C(N) stars. Unfortunately, the reproducibility of the TDU phenomenon is a main problem with AGB computations, connected to the more general astrophysical problem of the treatment of turbulent convection in stellar interiors. Different evolutionary codes do not reproduce the same results: some codes found a large amount of TDU at solar metallicity only for intermediate-mass stars \\citep[5 \\msb $\\lsim M \\lsim$ 8 \\ms,][]{ibena}, others found TDU to occur in low-mass stars ($M <$ 5 \\ms), but only for low metallicities: 1/200 of solar \\citep{ibenrenzinia} and 1/20 of solar \\citep{boothroydd}. \\citet{lattanzioc} and \\citet{stranieroc} found TDU to occur also in stars with initial mass as low as 1.5 \\msb and solar metallicity. The amount of TDU is also connected to the evolution of the envelope depending, for example, on the temperature at its base \\citep{wood} and its total mass \\citep{stranieroc}. The inclusion of extra mixing in the computation of TDU, however, can affect these results. The lower limit of 1.5 \\msb at solar metallicity for TDU to occur is consistent with inferences derived from observational constraints of carbon stars \\citep*{groenewegenvd}. \\citet{frost} studied the effect of the numerical treatment on the occurrence and magnitude of TDU, while \\citet{mowlavi} tested the effect of using an extra mixing mechanism at the base of the convective envelope and found that TDU occurs in his models only if this procedure is applied. Recently, \\citet{pols} showed that when the AGB evolution is computed by solving simultaneously the structure equations and the diffusion equations for changes in chemical composition the TDU strongly depends on numerical choices, such as how the diffusion coefficients are averaged.  Another debated problem is a satisfactory description of mass loss during the AGB phase. It is believed that the main mechanism responsible for mass loss in red giants is radiation pressure on dust grains. Detailed calculations in conjunction with new observations will address this problem. Most recent works include \\citet{wachter} and \\citet{olofsson}. In the past an often employed way to include mass loss has been to use \\citet{reimers}'s formula, in which the mass loss rate depends on the luminosity, radius and mass of the star and is proportional to a free parameter, $\\eta$, whose value can vary (typically from $\\simeq$ 0.3 to ten). Other semi-empirical relations for the mass-loss rate have been proposed, in particular by \\citet{vassiliadis,blocker} and \\citet*{arndt}. \\citet{marigo} investigated the effects of molecular opacities on the evolution of AGB stars as the C/O ratio grows from below to above unity and found that in the carbon-rich models a notable cooling results and hence an early onset of the {\\it superwind} is expected in the framework of the \\citet{vassiliadis}'s formulation of the mass loss rate.  \\subsection{The $s$ process}  The He intershell is the site of the $s$ process, the neutron capture nucleosynthesis that occurs when neutron number densities $N_{n}$ are of the order of $10^6 - 10^{11}$ n cm$^{-3}$. In this range of $N_{n}$ when an unstable nucleus is produced along the $s$-process chain of neutron captures, it typically decays rather than capture another neutron \\citep{burbidge}. However, for relatively high values of neutron density or temperature several $branchings$ along the $s$-process path are open and neutron-rich nuclei can be produced \\citep{ward}. As a ``rule of thumb'' the $s$ process distribution follows the law $\\sigma(A) N_s(A) = constant$, where $\\sigma(A)$ is the neutron capture cross section of nucleus $A$ and $N_s(A)$ its s-process abundance by number. This rule is only valid locally because of the presence of the neutron magic nuclei at $N$ = 50, 82, 126. These nuclei have very low neutron capture cross sections, which make them behave as bottlenecks for the neutron flux. As a consequence steps in the otherwise $\\sigma(A) N_s(A) = constant$ distribution are found in correspondence to neutron magic nuclei. The depth of these steps is a function of the total time-integrated neutron exposure $\\tau = \\int^t_0 N_n v_T dt$ (where $v_T$ is the thermal velocity), so that different neutron exposures lead to different $\\sigma(A) N_s(A)$ distributions \\citep[see][ page 568]{clayton}. In AGB stars heavy $s$-process elements are produced by the $s$ process and then mixed to the stellar surface by TDU where they have been observed \\citep[for general references see] []{smithlb,wallerstein,wallersteink,bussogl,abiab}. In particular, AGB stars are responsible for the $s$-process {\\it main component} \\citep{kappelerbw,arlandini,travagliogg}, i.e. a significant fraction of the Galactic production of elements from Sr to Pb, as well as for the {\\it strong component} \\citep{claytonr} feeding a large fraction of solar $^{208}$Pb \\citep{gallinoab,goriely,travagliogb}.  The \\ct($\\alpha$,n)$^{16}$O reaction is activated at low temperatures ($\\gsim 0.8 \\times 10^{8}$ K), and plays the major role as neutron source in AGB stars \\citep{hollowell,gallinobprr,kappelergb}. However, a higher amount of \\ctb than present in the H-burning ashes is needed to reproduce the observed enhancements of heavy elements. To enable the formation of a region rich in \\ct, the \\ctb {\\it pocket}, a few protons from the envelope must penetrate into the He intershell and then react with the \\cdb nuclei abundantly present there. The possibility of penetration of protons from the envelope directly into the convective PDCZ was found by \\citet{ibenc} to be severely inhibited by an entropy barrier induced by the still active H-burning shell. A favorable location for the mixing of protons is the sharp H/He discontinuity that is left over after the occurrence of TDU while the H shell is temporarily extinguished \\citep{ibenrenzinib}. Recently, models including time-dependent overshoot \\citep{herwigbs} motivated by hydrodynamic simulations by \\citet*{freytag} \\citep*[see also][]{singh}, and models with rotation \\citep{langer} have been able to create a proton rich layer with $M_{pocket} \\sim 10^{-5} - 10^{-6}$ \\msb at the top of the He intershell after the end of the TDU. A similar effect was obtained by imposing a velocity gradient profile beyond the border of the convective envelope \\citep{cristallo}. Another description of turbulent perturbation of stable layers below convective zones can be performed in the framework of internal gravity waves \\citep{denis}. In summary, the occurrence of hydrogen penetration at the top of the He intershell seems plausible, but the mass involved and the resulting \\ctb abundance profile are still to be treated as relatively free parameters, constrained by the rates of proton captures on the carbon isotopes.  The \\s-process mechanism in AGB stars is visualized in Fig. \\ref{fig1}. After less than a few thousand years \\ctb and \\nqb are produced by proton capture on the abundant \\cdb in the top layers of the He intershell. Subsequently, after about 10,000 yr, the \\ctanb reaction is activated during the interpulse phase in radiative conditions at a relatively low temperature: $\\sim 0.9 \\times 10^{8}$ K \\citep{stranierog}. Before the end of the interpulse period all \\ctb has burnt and the pocket has become enriched in \\s-processed material. This $s$-process nucleosynthesis occurs in radiative conditions, hence there is no interaction among the different layers of the pocket. Each zone can be treated separately and the efficiency of the neutron flux is a local property that only depends on the initial p/\\cdb ratio in that layer. The neutron flux lasts typically 20,000 yr and produces locally high neutron exposures, of the order of 10$^{-1}$ mbarn$^{-1}$. On the other hand, because the timescale is long, the neutron density remains low, reaching about $10^7$ n cm$^{-3}$. At the end of the interpulse the pocket is engulfed by the following PDCZ and mixed with ashes from the H-burning shell and material from the previous PDCZ. A large amount of \\nvdb is present in the PDCZ as a product of the chain \\nqagb starting on the abundant \\nqb from the H-burning ashes. In the convective zone the \\nvdanb reaction can be marginally activated and a second neutron flux occurs. This second neutron burst is opposite in features to that in the pocket: it occurs on a timescale of a few years and it produces low neutron exposures (of the order of $10^{-2}$ mbarn$^{-1}$) with a high-peaked neutron density, up to $10^{10}$ n cm$^{-3}$. This neutron burst does not typically contribute much to the overall production of $s$ elements, however it affects the final abundances of isotopes connected to branching points. For an extensive review on the $s$ process in AGB stars see \\citet*{bussogw}.  We want to analyze the dependence of \\s-process model predictions on the stellar evolutionary code used to compute them. We compare the results of post-processing models based on three evolutionary sequences computed with three different codes, one of them including an overshoot mechanism. The sequences are different in the details of the description of physical processes like mass loss, and of the numerical solution scheme adopted. To our purposes these computations represent a typical set of models of the same star constructed by different groups. The paper is organized as follows: in \\S 2 we describe the stellar evolutionary codes and the evolutionary sequences we used to compute the nucleosynthesis. In \\S 3 we describe the nucleosynthesis codes, the formation and efficiency of the \\ctb neutron source, and the neutron fluxes and final $s$-process distributions as computed with the different evolutionary sequences. In \\S 4 we compare the final results with a number of observational constraints. In \\S 5 we outline our conclusions. A preliminary study was presented by \\citet{lugaroh}. ",
        "conclusions": "We presented detailed computations of the $s$-process abundance distribution of heavy elements for a 3 \\msb star of solar metallicity making use of three different evolutionary codes. One of them includes a time-dependent hydrodynamic overshoot mechanism and finds the \\ctb pocket self-consistently. Hydrodynamic overshoot mechanism at the base of the convective envelope is a likely process to produce a \\ctb pocket. The alternative mixing process related to rotation may limit the neutron flux by continued mixing of \\nqb into the $s$-process layer during the interpulse phase \\citep{herwigll}. A combination of both processes should be investigated in the future. It is of course conceivable that other mixing processes could be involved, which may have complementary or similar properties to those of the hydrodynamic overshoot mechanism considered here. Overshoot at the base of the convective envelope generates higher TDU efficiencies making it possible to create carbon stars at very low luminosities. The inclusion of overshoot at the base of the PDCZ, and the consequent intershell dredge up, help to explain the observed carbon and oxygen abundances of  H-deficient central stars of planetary nebulae.  The two computations that do not include overshoot do not self-consistently find any proton diffusion at the end of TDU and we assumed an identical \\ctb pocket for these cases. The final results of these two cases are remarkably similar and we can conclude that the main uncertainty in \\s-process predictions is the \\ctb pocket and not the intrinsic differences among different codes when no overshoot mechanism is included. In the case that includes hydrodynamic overshoot, because of penetration of the intershell convective region into the CO core, (i) the abundance of \\cdb in the intershell and hence the maximum total neutron exposure achieved in the \\ctb pocket are higher than in the other cases, and (ii) the neutron flux in the PDCZs is stronger. In summary, overshoot at the base of the convective envelope produces a \\ctb pocket if large efficiencies, i.e. values of the parameter $f$, are assumed; overshoot at the base of the PDCZ influences the activation of the $^{13}$C($\\alpha$,n)$^{16}$O and $^{22}$Ne($\\alpha$,n)$^{25}$Mg neutron sources producing a more efficient neutron flux in both situations. The final resulting abundance distributions of heavy elements reflect such features of the neutron fluxes: in the case with overshoot included, the heaviest elements, such as Sm, are produced in greater amount than lighter elements such as Sn. At the end of the $s$-process path $^{208}$Pb is produced in a higher amount, as are also isotopes fed by branching points, such as $^{86}$Kr, $^{87}$Rb and $^{96}$Zr.  With our choice of the \\ct-pocket profile the results from the two cases without overshoot give a good match to several observational constraints. The strong intershell dredge up connected with overshoot at the base of the PDCZs is in several situations in contrast with observational constraints. We still need to expand this study to a range of masses and metallicities. Within the EVOL case, we also need to extend our calculations in the space of the free overshoot parameter $f$, as applied at different boundaries and possibly also different times during the evolution, in order to test which range of $f$ could provide us with the closest match to observational constraints. For example, a preliminary test with $f$=0.008 applied at the base of the PDCZ show a significantly reduced efficiency in the activation of the $^{22}$Ne neutron source."
    },
    "0212/astro-ph0212128_arXiv.txt": {
        "abstract": "We describe a general yet simple method to analyse the propagation of nuclear reaction rate uncertainties in a stellar nucleosynthesis and mixing context. The method combines post-processing nucleosynthesis and mixing calculations with a Monte Carlo scheme. With this approach we reanalyze the dependence of theoretical oxygen isotopic ratio predictions in first dredge-up red giant branch stars in a systematic way. Such predictions are important to the interpretation of pre-solar Al$_{\\rm 2}$O$_{\\rm 3}$ grains from meteorites. The reaction rates with uncertainties were taken from the NACRE compilation \\citep{nacre}. We include seven reaction rates in our systematic analysis of stellar models with initial masses from 1 to 3$~{\\rm M}_{\\sun}$. We find that the uncertainty of the reaction rate for reaction ${\\rm ^{18}O(p,\\alpha)^{15}N}$ typically causes an error in the theoretical ${\\rm ^{16}O/^{18}O}$ ratio of $\\simeq +20 /$$-5$ per cent. The error of the  ${\\rm ^{16}O/^{17}O}$ prediction is 10--40 per cent depending on the stellar mass, and is persistently dominated by the comparatively small uncertainty of the  ${\\rm ^{16}O(p,\\gamma)^{17}F}$ reaction. With the new estimates on reaction rate uncertainties by the NACRE compilation, the p-capture reactions ${\\rm ^{17}O(p,\\alpha)^{14}N}$ and ${\\rm ^{17}O(p,\\gamma)^{18}F}$ have virtually no impact on theoretical predictions for stellar mass $\\leq 1.5~{\\rm M}_{\\sun}$. However, the uncertainty in ${\\rm ^{17}O(p,\\alpha)^{14}N}$ has an effect comparable to or greater than that of ${\\rm ^{16}O(p,\\gamma)^{17}F}$ for masses $> 1.5~{\\rm M}_{\\sun}$, where core mixing and subsequent envelope mixing interact. In these cases where core mixing complicates post-dredge-up surface abundances, uncertainty in other reactions have a secondary but noticeable effect on surface abundances. ",
        "introduction": "Observations of abundances and abundance ratios in stars yield powerful constraints on models of internal processes in stars. In particular, the predicted abundance evolution of stars of a given mass and metallicity depends on both mixing and nuclear burning processes. While it is customary to provide observational results with some estimate or analysis of the associated errors, this is seldomly done for theoretical abundance predictions. This is partly due to the fact that quantitative uncertainties on the input physics is not available and that theoretical model uncertainties (in particular for poorly understood mixing processes) are hard to make. Finally, a systematic error propagation in stellar models can be computationally demanding and the effort has often not been justified in the past. However, new spectroscopic observations with unprecedented precision will be available in the near future, and one has the impression that stellar physics in general is entering a new high-precision era. In this context it seems necessary to reconsider the question of abundance prediction uncertainties in a quantitative way.  As an initial step in this direction, we have chosen specifically the dependence of oxygen isotopic ratio predictions in giant stars on nuclear reaction rate uncertainties. The oxygen isotopic ratios are of particular interest because they are affected not only by nucleosynthesis but also by mixing processes which are not yet very well understood. The latter is evident from the spectroscopic properties of giant stars which are not entirely consistent with standard stellar evolution models \\citep{harris85,harris:87,boothroyd:99}. Therefore it is important to quantitatively know the uncertainties arising from nucleosynthesis in order to help constrain mixing processes using observations. In addition, the oxygen isotope predictions are important for the interpretation of pre-solar corundum (Al$_{\\rm 2}$O$_{\\rm 3}$) from meteorites for which oxygen isotopic ratios can be measured with high precision \\citep[e.g.][]{huss,nittler97}.  During the main sequence (MS) evolution of low- and intermediate-mass stars, partial H-burning in the envelope produces $^{17}$O, $^{13}$C, and $^{14}$N while $^{12}$C, $^{15}$N, and $^{18}$O are destroyed (Fig.\\,\\ref{fig:abund}). The abundance of $^{16}$O in the envelope is essentially unchanged for stellar mass $\\leq 1.5~{\\rm M}_{\\sun}$. For more massive stars, core convection on the early MS reaches out to $\\sim 25$ per cent of the stellar mass, leaving behind core processed material as it retreats, including destroyed $^{16}$O \\citep[e.g.][]{schaller:92}. Evidence of the retreating core convection is seen in the 2 and 3$~{\\rm M}_{\\sun}$ cases in Fig.\\,\\ref{fig:abund} at the bottom of the envelope where the H and $^{16}$O are depleted.  At the end of the MS phase hydrogen is exhausted in the core, causing the core to contract and the envelope to expand and cool. Envelope convection descends into the star as the envelope temperature drops. Regions previously affected by cool H-burning are homogenised, leading to a change of surface abundances (1$^{st}$ dredge-up, 1dup hereafter). As a result of this mixing, theory predicts the $^{16}$O/$^{17}$O ratio decrease from initially $\\sim 2600$ to a few hundred, depending on stellar mass while the $^{16}$O/$^{18}$O ratio increases only marginally by $\\leq 20$ per cent from the initial value of about 500, weakly dependent on mass. For many stars, this is in fair agreement with spectroscopic observations. For example, \\citet{harris84} find for K-giants ${\\rm ^{16}O/^{17}O}$ ratios in the range $300 \\dots 1000$ and ${\\rm ^{16}O/^{18}O} \\sim 425 \\dots 600$. Similarly, some corundum oxygen isotopes are in rough agreement with standard 1dup model predictions (e.g.\\ SEAL203 in \\citet{choi:98} has  ${\\rm^{16}O/^{17}O}=355$ and a solar   ${\\rm ^{16}O/^{18}O}$ ratio) while many others require either non-solar initial ${\\rm ^{16}O/^{18}O}$ ratios \\citep{huss,boothroyd94,timmes95} or non-standard mixing processes.  The dependence of oxygen isotopic ratios on the ${\\rm^{17}O}$ proton capture nuclear reactions has been studied before by \\citet{eleid} and \\citet{boothroyd94}. They considered the effect of a new determination of these reactions by \\citet{landre:89} and found that the uncertainty of the  ${\\rm ^{16}O/^{17}O}$ predictions is dominated by the uncertainty of these rates for stars with initial masses larger than $\\sim 2.5{\\rm M_{\\sun}}$.  In this paper we include seven reactions in a systematic analysis of the influence of rate uncertainties  on the predicted oxygen isotopic ratios in red giant branch (RGB) stars. In addition to addressing this scientific question, we also want to demonstrate and test the method of evaluating model error propagation described here. Sect. 2 describes the method and computations. In Sect. 3 the uncertainty in the modeled oxygen isotopic ratios of RGB stars is analysed for significant reaction rate uncertainties and comparisons are made to observations and other models. A summary is in the final section. ",
        "conclusions": "In this paper we quantify the uncertainty in oxygen isotopic ratios due to uncertain reaction rates. These results may help to motivate and prioritize new laboratory measuremtents of reaction rates. Reaction rate uncertainty for ${\\rm ^{16}O(p,\\gamma)^{17}F}$ and ${\\rm ^{18}O(p,\\alpha)^{15}N}$ are significant to oxygen isotope ratios for all of the stellar cases studies here. Reaction ${\\rm ^{17}O(p,\\alpha)^{14}N}$ competes with ${\\rm ^{16}O(p,\\gamma)^{17}F}$ for dominance in $^{16}$O/$^{17}$O uncertainty for the 2--3$~{\\rm M}_{\\sun}$ cases because of the interaction between core mixing and subsequent envelope mixing. Of course, this result is based on the assumption that the estimated uncertaities in the reaction rates (NACRE upper and lower limits) are appropriate.  In general, reactions with large uncertainties, and reactions with slow reaction rates but large production/destruction rates and {\\it any} uncertainty create the largest uncertainty in isotopic ratios. The MC scheme demonstrated here efficiently finds the most problematic rates and provides a means of quantifying uncertainties for a particular stellar environment.  The model uncertainties calculated here are from reaction rate uncertainties only. The total model uncertainty in oxygen istopic ratios definitly has contributions from uncertainty in mixing processes. This problem will become more tractable as the uncertainties due to reaction rates are made smaller."
    },
    "0212/astro-ph0212291_arXiv.txt": {
        "abstract": "We present the results of time-dependent, numerical magnetohydrodynamic simulations of a realistic young stellar object outflow model with the addition of a disk-associated magnetic field. The outflow produced by the magnetic star-disk interaction consists of an episodic jet plus a wide-angle wind with an outflow speed comparable to that of the jet (100--200 km s$^{-1}$).  An initially vertical field of $\\ll 0.1$ Gauss, embedded in the disk, has little effect on the wind launching mechanism, but we show that it collimates the entire flow (jet + wide wind) at large (several AU) distances. The collimation does not depend on the polarity of the vertical field. We also discuss the possible origin of the disk-associated field. ",
        "introduction": "Prevalent theoretical models for winds launched from accretion disks (see, e.g., \\opencite{koniglpudritz00} for a review) hold that the final wind velocity is of the order of the Keplerian rotational velocity of the launch point.  The high velocity of observed jets from young stellar objects (YSO's) suggests that they originate from a deep potential well \\cite{kwantademaru88}, requiring the launching region to be less than an AU in extent, for reasonable parameters.  We therefore adopt the view that these jets are launched from a region within several stellar radii.  The observations require that the flows are launched initially with large opening angles and become collimated within a few 10's of AU \\cite{eisloffelea00}.  Since the winds become collimated along the rotational axis of the accretion disk, the collimation process must be associated with the disk.  In this work, we explore one possible explanation, that disk-associated, poloidal magnetic fields collimate a central, fast, wide-angle wind into an optical jet. ",
        "conclusions": ""
    },
    "0212/astro-ph0212258_arXiv.txt": {
        "abstract": "The AF-supergiants in the Galaxy and the SMC allow us to examine predictions from evolution models through their CNO abundances. In these proceedings, we recalculate the NLTE nitrogen abundances in 22 Galactic and 9 SMC A-supergiants using improved atomic data and model atmospheres to compare with new evolution models. The new abundances are higher than previously published values, and suggest that most of these stars have undergone substantial mixing with CN-cycled gas.   While there is no clear relationship with mass, there is an apparent relation with metallicity since the SMC stars (including B-stars) have larger nitrogen enrichments. We suggest that rotational mixing is indicated from the main-sequence throughout the supergiant range, with more substantial rotational mixing in the SMC stars.   In addition, the SMC AF-supergiants appear to have undergone the first dredge-up during a previous red giant phase, and possibly the Galactic AF-supergiants have as well.   All abundances are compared to the new solar abundances from M. Asplund (this conference). ",
        "introduction": "Reviews on massive star evolution have shown there are a number of observations that suggest an additional parameter affects the evolution scenarios, ranging from blue-to-red supergiant ratios in clusters to CNO abundances in main-sequence and evolved stars; c.f., Maeder \\& Meynet (2000), Heger \\& Langer (2000 = HL00), Maeder \\& Conti (1994). In Galactic supergiants, the N/C ratios have long supported a parameter that varies from star-to-star, like rotation.   This is because a variety of N/C ratios have been found within clusters of stars (thus, not natal variations), and because the N/C ratios do not clearly scale with effective temperature, luminosity, or mass.  Evolution models that include effects of rotation emerged in 2000. Not only can rotation help to explain the abundance ratios and certain mass discrepancies, but the new models also predict that massive stars may produce and mix primary nitrogen into their atmospheres (by Meynet \\& Maeder 2002).    The new models have been calibrated based on abundance ratios, masses, rotation rates, and ages in the published literature.  But they also make definate predictions that can be tested for further refinement.  Unfortunately, one direct test that is difficult to implement is to examine abundance ratios versus rotation rate.   Rarely do we know the inclination angle of a rotating star.   Also, as $v$sin$i$ increases, then the spectral lines become so shallow and broad that the atmospheric analysis becomes quite difficult and less certain.   Fortunately, there are other indirect tests. ",
        "conclusions": ""
    },
    "0212/astro-ph0212550_arXiv.txt": {
        "abstract": "The search for Earth-like extrasolar planets is in part motivated by the potential detection of spectroscopic biomarkers. Spectroscopic biomarkers are spectral features that are either consistent with life, indicative of habitability, or provide clues to a planet's habitability. Most attention so far has been given to atmospheric biomarkers, gases such as O$_2$, O$_3$, H$_2$O, CO, and CH$_4$.  Here we discuss surface biomarkers. Surface biomarkers that have large, distinct, abrupt changes in their spectra may be detectable in an extrasolar planet's spectrum at wavelengths that penetrate to the planetary surface. Earth has such a surface biomarker: the vegetation ``red edge'' spectroscopic feature. Recent interest in Earth's surface biomarker has motivated Earthshine observations of the spatially unresolved Earth and two recent studies may have detected the vegetation red edge feature in Earth's hemispherically integrated spectrum. A photometric time series in different colors should help in detecting unusual surface features in extrasolar Earth-like planets. ",
        "introduction": "One hundred extrasolar giant planets are currently known to orbit nearby sun-like stars.  These planets have been detected by the radial velocity method and so, with the exception of the one transiting planet, only the minimum mass and orbital parameters are known.  Many plans are underway to learn more about extrasolar planets' physical properties from ground-based and space-based observations and via proposed or planned space missions.  Direct detection of scattered or thermally emitted light from the planet itself is the only way to learn about a variety of the planet's physical characteristics. Direct detection of Earth-size planets, however, is extremely difficult because of the proximity of a parent star that is $10^6$ to $10^{10}$ times brighter than the planet.  Terrestrial Planet Finder (TPF), with a launch date in the 2015 timeframe, is being planned by NASA to find and characterize terrestrial-like planets in the habitable zones of nearby stars. The ESA mission Darwin has similar goals. The motivation for both of these space missions is the detection and spectroscopic characterization of extrasolar terrestrial planet atmospheres. Of special interest are atmospheric biomarkers---such as O$_2$, O$_3$, H$_2$O, CO and CH$_4$---which are either indicative of life as we know it, essential to life, or can provide clues to a planet's habitability (\\cite[Des Marais \\etal\\ 2002]{DM02}).  In addition, physical characteristics such as temperature and planetary radius could be constrained from low-resolution spectra.  We have shown (\\cite[Ford, Seager \\& Turner 2001]{FST01}) that planet characteristics could also be derived from photometric measurements of the planet's variability at visible wavelengths. A time series of photometric data of a spatially unresolved Earth-like planet could reveal a wealth of information such as weather, the planet's rotation rate, presence of large oceans or surface ice, and existence of seasons. The amplitude variation of the time series depends on cloud-cover fraction; more cloud cover makes a more photometrically uniform Earth and so reduces variability. The signal-to-noise necessary for photometric study would be obtained by a mission capable of measuring the sought-after atmospheric biomarker spectral features. Furthermore the photometric variability could be monitored concurrently with a spectroscopic investigation, as was done for the transiting extrasolar giant planet HD209458b (\\cite[Charbonneau \\etal\\ 2002]{Charb02}).  To detect and study surface properties only wavelengths that penetrate to the planetary surface are useful. Visible wavelengths are more suited than mid-IR wavelengths for such measurements for several reasons. First, the albedo contrast of surface components is much greater than the temperature variation across the planet's surface. Second, at visible wavelengths the planet's flux is from scattered starlight and hence at some configurations the planet is only partially illuminated. This allows a more concentrated signal from surface features, such as continents, as they rotate in and out of view. Furthermore, the non-uniform illumination and non-isotropic scattering of different surface components mean much of the scattered light can come from a small part of the planet's surface. At mid-IR wavelengths the planet has, to first order, uniform flux across the planet hemisphere. In addition, the narrow transparent spectral window at 8-12~$\\mu$m will close for warmer planets than Earth and for planets with more water vapor than Earth. However, further study at the mid-IR ``window'' needs to be investigated.  An extremely exciting possibility, aided by a photometric time series, is the detection of surface biomarkers in the spectrum of an extrasolar planet. This would be possible at wavelengths that penetrate to the planet's surface, and for surface features that have large, distinct, abrupt changes in their spectra. Although most surface features (e.g., ice, sand) show very little or very smooth continuous opacity changes with wavelength, Earth has one surface feature with a large and abrupt change: vegetation (Figure~\\ref{fig:plantspectrum}). In this paper we discuss Earth's vegetation red-edge spectroscopic feature as a surface biomarker. ",
        "conclusions": "The vegetation red edge spectroscopic feature is a factor of 5 or more change in reflection at $\\sim$ 0.7$\\mu$m. This red edge feature is well-used in satellite remote sensing studies of Earth's vegetation. Earthshine observations have been used to detect the vegetation red edge signature in the spatially unresolved spectrum of Earth where it appears at the few percent level.  When discovered, observations of extrasolar Earth-like planets at wavelenghts that penetrate to the planet's surface will be very useful, especially for planets with much lower cloud cover than Earth's 50\\%.  A time series of spectra or broad-band photometry could reveal surface features of a spatially unresolved planet, including surface biomarkers. Earth's hemispherically integrated vegetation red-edge signature is weak (a few to ten percent), but Earth-like planets with different rotation rates, obliquities, land-ocean fraction, and continental arrangement may well have lower cloud-cover.  While it is near impossible to speculate on spectral features of light harvesting organisms on extrasolar planets, flexible data acquistion will maximize scientific return. The detection of an unusual spectral signature that is inconsistent with any known atomic, molecular, or mineralogical signature would be fantastic. Combinations of unusual spectral features together with strong disequilibrium chemistry would be even more intriguing and would certainly motivate additional studies to better understand the propsects for such a planet to harbor life."
    },
    "0212/astro-ph0212416_arXiv.txt": {
        "abstract": "As part of the Hubble Deep Field South program, a set of shorter 2-orbit observations were obtained of the area adjacent to the deep fields. The WFPC2 flanking fields cover a contiguous solid angle of 48 square arcminutes. Parallel observations with the STIS and NICMOS instruments produce a patchwork of additional fields with optical and near-infrared (1.6 $\\mu m$) response. Deeper parallel exposures with WFPC2 and NICMOS were obtained when STIS observed the NICMOS deep field. These deeper fields are offset from the rest, and an extended low surface brightness object is visible in the deeper WFPC2 flanking field. In this data paper, which serves as an archival record of the project, we discuss the observations and data reduction, and present SExtractor source catalogs and number counts derived from the data. Number counts are broadly consistent with previous surveys from both ground and space. Among other things, these flanking field observations are useful for defining slit masks for spectroscopic follow-up over a wider area around the deep fields, for studying large-scale structure that extends beyond the deep fields, for future supernova searches, and for number counts and morphological studies, but their ultimate utility will be defined by the astronomical community. ",
        "introduction": "\\label{s:introduction}  The history of the Hubble Deep Field-North (HDF-N, Williams et al.~1996) represents an example of the benefits obtained from technological advances and observing from space. Following the first HST servicing mission, relatively deep observations of the cluster CL0939+4713 at z=~0.4 (Dressler et al.~1994) and the cluster around 3C324 at z=~1.2 (Dickinson et al.~1995), as well as observations of field galaxies from the HST Medium Deep Survey (Griffiths et al.~1996) demonstrated that the Hubble Space Telescope with its restored optical capability could yield images of distant galaxies with unprecedented clarity, revealing many faint galaxies and a wealth of morphological detail in objects observed. The Hubble Deep Field-North was the culmination of these early efforts.  In order to provide a similar opportunity for follow-up by observatories in the southern hemisphere, and to provide comparison along a different line of sight another deep-field campaign was carried out in October, 1998. For this second field, a line-of-sight to a QSO at redshift $z=2.2$ was chosen, offering the opportunity to correlate properties of the intergalactic medium to the density of surrounding galaxies. New instruments on HST also provided the opportunity for ultraviolet and infrared imaging on the second campaign, as well as providing spectroscopy and deep unfiltered optical imaging. Scientific results from the Hubble Deep Fields are reviewed by Ferguson, Dickinson, \\& Williams (2000). Further details of the HDF-S observations can be found in Williams et al. (2000), Casertano et al. (2000), Gardner et al. (2000), and Fruchter et al (2002).  The Hubble Deep Field-North (HDF-N, Williams et al.~1996) program included a set of shallower Wide Field Planetary Camera 2 (WFPC2) observations immediately adjacent to the deep WFPC2 field.  These fields were included for the benefit of observers who could use the objects in those fields as additional targets when obtaining spectroscopic observations of the primary objects in the Deep Field itself (e.g. Cohen et al. 2000). The HDF-N flanking fields (FFs) have proven to be scientifically useful in a variety of ways. They have provided morphological and photometric data for targets of ground-based spectroscopy. They have contributed to studies of Galactic structure via star counts, as well as studies of extragalactic structures on a scale larger than the Deep Field, and have provided detections and morphologies of objects identified in X-ray, radio, infrared, and sub-mm surveys.  Similar flanking field observations were included in the Cycle 7 Hubble Deep Field-South (HDF-S, Williams et al.~2000) campaign which included the newly-available Space Telescope Imaging Spectrograph (STIS) and the Near-Infrared Camera Multi-Object Spectrometer (NICMOS) as well as WFPC2. The quasar J2233-606 was centered in the STIS field and observed both spectroscopically and in imaging mode at a variety of wavelengths. NICMOS and WFPC2 provided simultaneous multi-color imaging in nearby parallel deep fields. The availability of STIS and NICMOS in addition to WFPC2 has augmented the scientific utility of the HDF-S FFs, allowing us not only to probe the extended environment of the QSO in the area between the deep fields, but also other ``core samples'' at somewhat greater distances.  The HDF-S flanking-fields campaign consisted of 27 orbits of observations. Nine contiguous WFPC2 fields were observed for two orbits each, and exposures were taken in parallel with NICMOS and STIS. The remaining nine orbits were invested in a STIS clear-filter image of the NICMOS deep field (STIS-on-NICMOS), accompanied by parallel observations with the other two instruments. This campaign resulted in shallow WFPC2 coverage of a contiguous area stretching between the STIS and NICMOS deep fields, surrounded by a set of offset fields with coverage by a single instrument. The resulting patchwork of observations is illustrated in Figure 1 of Williams et al. (2000).  A version of the STIS-on-NICMOS image which was convolved with the NICMOS PSF, and the catalog derived from that image, will be presented in the paper by Fruchter et al. (2002) on the NICMOS deep field. The unconvolved STIS-on-NICMOS image and the catalog derived from it, which has a somewhat fainter magnitude limit than that derived from the NICMOS PSF-convolved image, will be presented in this paper.  Williams et al. (2000) gives a general overview of the HST observations of the HDF-S, and describe the rationale for adopting the field around QSO J2233-606 for the campaign. The present paper describes the HDF-S flanking field observations and data processing, and presents source catalogs and number counts. In Section 2, we describe the observations, while Sections 3, 4, and 5 describe data reduction and image combination for WFPC2, STIS, and NICMOS FFs, respectively. Astrometry updates and large mosaiced images are covered in Section 6, and in Sections 7 and 8 we present the catalogs and number counts. ",
        "conclusions": "\\label{s:conclusions}  We have presented the entire suite of HDF-South flanking field observations in all three instruments, describing the observations, and data reduction processes. We have reviewed updates to the astrometry and have described extended mosaics of multiple flanking fields in detail. We have also presented catalogs and number counts derived from them. The number counts are broadly consistent with those of previous surveys. The original driving force for the flanking fields was the need for more targets beyond the deep fields for long-slit spectroscopy at ground-based telescopes. In meeting this need, this large dataset is also useful for other purposes, especially studies of galaxy evolution involving number counts, sizes, and morphologies, as well as studies of potential absorbers along the lines of sight near the QSO. They may also be used as ``baseline'' images for later follow-up observations to hunt for distant supernovae. These data can also be useful for galactic studies such as star counts. While their utility has mainly been seen by us as involving the kinds of studies mentioned above, there are undoubtedly many other uses of which we have not yet conceived, and which may represent their greatest potential."
    },
    "0212/astro-ph0212370_arXiv.txt": {
        "abstract": "The power of blazar jets rivals the power that gravity can extract from accreting matter. The mechanism launching and accelerating jets can be considered as the most efficient engine operating in radio--loud sources. It is still a matter of debate if the jet carries this power to the radio lobes, hundreds of kpc away, in the form of Poynting flux or bulk kinetic energy, or both, and if these two ingredients have relative weights changing along the way. Accordingly, there are two (or more) possible general scenarios for how the jet can dissipate part of its power into radiation. It can be through e.g. reconnection of the magnetic field in the purely electromagnetic scenario, or through internal shocks in the matter dominated picture. Ways to discriminate these ideas are welcome. ",
        "introduction": "The bulk of the radiation we observe from blazars is produced in a well localized region of their relativistic jets. This is, in my opinion, one of the most important results of the $\\gamma$--ray observations performed by EGRET onboard the Compton Gamma Ray Observatory, and of the more recent observations performed by ground based Cherenkov telescopes. When observational campaigns were organized at $\\gamma$--ray and other frequency bands, they proved that also the X--ray flux was varying at the same time of the $\\gamma$--ray flux, even if with a smaller amplitude. By demonstrating that both fluxes are likely to be produced in the same region, this puts a strong limit on the local photon energy densities, which must be sufficiently small to avoid the absorption caused by the photon--photon pair production process. This implies that the observed flux is enhanced by beaming and therefore that the source is moving relativistically. The same transparency constraint requires also that the $\\gamma$--ray emitting region of the jet is sufficiently far from the X--ray emitting corona sandwiching the accretion disk. On the other hand, the hour--day variability timescales poses an upper limit on the size of the emitting region. The jet region where most of the dissipation takes place must then be at a few hundreds of Schwartzchild radii from the black hole (Ghisellini \\& Madau 1996)  The other important result of the high energy data is that we now know the entire spectral energy distribution (SED) of blazars: since the $\\gamma$--ray luminosity is often dominating the total radiation output, we now know the total radiated power. This is large. The total power in radiation only, accounting for beaming, equals $L_{\\rm obs}/\\Gamma^2$, where $L_{\\rm obs}$ is the bolometric luminosity calculated assuming isotropy and $\\Gamma$ is the bulk Lorentz factor. Jets must carry a power larger than that, implying for instance that the jet of S5 0836+71 (which has $L_{\\rm obs}\\sim 10^{49}$ erg s$^{-1}$, Tavecchio et al. 2000), should carry more than $10^{47}/(\\Gamma/10)^2$ erg s$^{-1}$. The question is in what form. It can be purely electromagnetic, as advocated by Blandford (2002) and by Lyutikov \\& Blandford (2002), or largely in the form of bulk motion of matter, with a relatively weak Poynting flux. Also, the relevant ``matter\" can be protons, or instead mostly electron--positron pairs. And of course there can be all the combinations of the above ingredients, possibly changing weight at different distances from the black hole: this underlines the importance to find ways to measure the power of jets and their matter content at different scales. For instance, one can calculate the amount of the emitting particles and the magnetic field required to account for the radiation we see, and if also the size and the bulk Lorentz factor are known, one can set a lower limit to the jet power (it is a lower limit because there can be other cold, non emitting, particles). This can be done at the jet scale where most of the radiation is produced (i.e. $10^{17}$ cm from the black hole, see e.g. Ghisellini 1999 and Fig. 1), where superluminal motion is detected (i.e. at the parsec--scale, see e.g. Celotti \\& Fabian 1991); at the large jet scales (i.e. at hundreds of kpc, as recently observed by Chandra, VLA, HST, see e.g. Ghisellini \\& Celotti 2001; Tavecchio et al. 2000); and finally at the scale of the radio lobes, thought as calorimeters (e.g. Rawlings \\& Saunders 1991). These studies point to the important conclusion that the engine generating the jet can be even more efficient than accretion. This would become dramatically clear if also gamma--ray bursts (GRBs) are jetted sources. ",
        "conclusions": ""
    },
    "0212/astro-ph0212146_arXiv.txt": {
        "abstract": "I compare the density profile of dark matter (DM) halos in cold dark matter (CDM) N-body simulations with 1 \\Mpc, 32 \\Mpc, 256 \\Mpc and 1024 \\Mpc box sizes. I compare the profiles when the most massive halos are composed of about $10^5$ DM particles. The DM density profiles of the halos in the 1 \\Mpc box at redshift $z \\sim 10$ show systematically shallower cores with respect to the corresponding halos in the 32 \\Mpc simulation at $z \\sim 3$ that have masses, $M_{dm}$, typical of the Milky Way and are fitted by a NFW profile. The DM density profiles of the halos in the 256 \\Mpc box at $z \\sim 0$ are consistent with having steeper cores than the corresponding halos in the 32 \\Mpc simulation, but higher mass resolution simulations are needed to strengthen this result.  Combined, these results suggest that the density profile of DM halos is not universal, presenting shallower cores in dwarf galaxies and steeper cores in clusters. More work is needed to validate this finding at $z=0$.  Physically the result sustains the hypothesis that the mass function of the accreting satellites determines the inner slope of the DM profile. But the result can also be interpreted as a trend with the dynamical state in the assembly process of halos of different mass.  In comoving coordinates, $r$, the profile \\[ \\rho_{dm} \\propto {1 \\over X^\\alpha(1+X)^{3-\\alpha}} \\] with $X=c_\\Delta r/r_\\Delta$ and $\\alpha =(9+3n)/(5+n) \\approx 1.3+(M_{dm,14}^{1/6}-1)/(M_{dm,14}^{1/6}+1)$, provides a good fit to all the DM halos from dwarf galaxies to clusters at any redshift with the same concentration parameter $c_\\Delta \\sim 7$. Here, $r_\\Delta(M_{dm})$ is the virial radius, $n$ is the effective spectral index of the initial power spectrum of density perturbations and $M_{dm,14}=M_{dm}/(3 \\times 10^{14}$ M$_\\odot$). The slope, $\\gamma$, of the outer parts of the halo appears to depend on the acceleration of the universe: when the scale parameter is $a=(1+z)^{-1} \\simlt 1$, the slope is $\\gamma \\approx 3$ as in the NFW profile, but $\\gamma \\approx 4$ at $a >1$ when $\\Omega_\\Lambda \\sim 1$ and the universe is inflating. The shape of the DM profiles presents a significant scatter around the mean. It is therefore important to analyse a significant statistical sample of halos in order to determine the mean profile.  I compare the DM profiles in the 1 \\Mpc box with the same simulation including stars, baryons and radiative transfer presented by Ricotti, Gnedin and Shull. Radiative feedback effects produce a larger scatter in the density profile shapes but, on average, do not affect the shape of the DM profiles significantly. ",
        "introduction": "According to the currently favoured galaxy formation scenarios, galaxies are formed from the gravitational growth of tiny density perturbations imprinted on a uniform universe during inflation. Small mass perturbations grow faster and constitute the building blocks for the assembly of larger galaxies and clusters. In the nonlinear phase of the gravitational collapse the dark matter (DM) particles are shock heated to the virial temperature and settle into a dark halo with mean overdensity about $\\Delta=178$ (according to the simple spherical collapse model in a flat $\\Omega_0=1$ universe).  Analysing N-body simulations of hierarchical structure formation \\cite*{Navarro:96,Navarro:97} (hereafter, NFW) have proposed that DM halos develop a universal density profile, valid for virialized halos at any redshift and with any mass, from dwarf galaxies to clusters. Further works have shown that an approximatively universal profile develops regardless of the details of the adopted cosmology and initial power spectrum of density perturbations \\citep{Huss:99,Eke:01}. The NFW density profile has a core cusp $\\rho \\propto r^{-1}$ and in the outer regions the density decreases as $\\rho \\propto r^{-3}$. The details of the gravitational collapse affect only the relative size of the core with respect to the virial radius (the halo concentration parameter, $c_\\Delta$) but not the slope of the profile. This result implies that during the virialization process the memory of the mass function and profile shapes of the building blocks that formed the halo is lost.  Recent observations of dwarf galaxies (spheroidal and irregular) and low surface brightness (LSB) galaxies seem to show that a flat core for the DM density profile is more compatible with the measurements of the rotation curves \\citep{vandenBosch:00,deBlok:01,deBlok:02, BorrielloS:01,SalucciB:00}.  These observations have renewed interest in the subject and a number of solutions have been proposed to solve this possible discrepancy.  The proposed solutions range from feedback effects of the baryons and stars on the DM density profile to modifications of the physical properties (\\eg, scattering cross section, temperature) of the DM particles. Some examples of alternatives to CDM are warm dark matter \\citep{BodeO:01} and self-interacting DM \\citep{Spergel:00}. The potential discrepancy between the steep DM cores found in N-body simulations with the observed flat core of dwarf galaxies is perhaps the most serious problem faced in CDM cosmologies today.  High resolution simulations and careful convergence tests have been done or are underway to understand if the discrepancies found by some authors (\\eg, \\cite{Moore:98,Moore:99} find steeper core profiles) are due to numerical artifacts.  In the core of DM halos the particle crossing time is much shorter than the Hubble time and the integration of trajectories is subject to subtle numerical effects that are difficult to keep under control.  \\cite{Power:02} estimated that about 1 million particles are needed in order to have convergent results in the inner 1\\% of the halo virial radius.  To have such a high number of particles per halo the use of a special technique is required. This technique consists of simulating with high mass resolution only the particles that will end up in the halo of interest at $z=0$, while having less mass resolution for the particles that end up far away from the halo of interest. This has the drawback that each simulation can resolve only one halo at a time, therefore selection effect biases and a poor statistical sample could affect the reliability of the final result even if the simulation is very accurate and has high resolution.  Note that most of the detailed and computationally expensive work on the subject has focused on simulating halos at $z=0$ with typical mass similar to the Milky Way.  The mass function of the building blocks of dwarf galaxies is different in comparison to the one for larger galaxies or clusters. On small scales (or masses) the mass fraction of virialized halos is about constant as a function of the logarithmic DM mass of the halos because the power spectrum has a slope $n \\sim -3$.  This means that small and large mass satellites contribute equally to the mass of the accreting halo.  In massive halos as in the Milky Way or in clusters of galaxies, the contribution from large mass accreting satellites is instead dominant.  For this reason it is more crucial to resolve very small mass satellites (\\ie, to have high mass resolution) in simulating dwarf galaxies than in simulating massive galaxies or clusters.  In this paper I compare the density profiles of DM halos in four CDM N-body simulations that are identical apart from the box sizes that are 1 \\Mpc, 32 \\Mpc, 256 \\Mpc and 1024 \\Mpc.  The simulation method is a two-level particle-particle particle-mesh (P$^3$M) with $256^3$ particles and $\\Lambda$CDM cosmology.  In dimensionless units the only difference between the simulations is the initial power spectrum of density perturbations. I compare the profiles when the most massive halos are composed of about $10^5$ DM particles and the clustering in the simulations is similar (see \\fig~\\ref{fig:hal}). This is done in order to minimise possible systematic errors in the simulations and to be more confident that any variation in the profile shapes is produced by the different initial power spectrum.  The masses of the halos studied in the four simulations, from the small to the large box, are typical of dwarf galaxies, Milky Way mass galaxies, clusters and superclusters of galaxies, respectively.  \\begin{figure} \\centerline{\\psfig{figure=pk.eps,width=9cm}} \\caption{\\label{fig:pk} Linear power spectrum of density fluctuations at $z=0$ (solid line) and logarithmic slope, $n(k)$, of the power spectrum (dashed line). The horizontal bars show the range of wave numbers, $k=2 \\pi/\\lambda$, resolved in the $1, 32, 256$ and $1024$ \\Mpc simulations.} \\end{figure}  If the cores of small mass halos are dense enough to resist disruption and survive ``undigested'' when incorporated into a larger object, they should determine the halo structure in the inner regions. Using scaling arguments \\cite*{SubramanianCO:00} have shown that, in a flat universe with an initial power spectrum of density perturbation $P(k) \\propto k^n$, the core density profile is $\\rho_{dm} \\propto r^\\alpha$ with $\\alpha = (9+3n)/(5+n)$. The numerical results that they report approximately confirm this expectation. In $\\Lambda$CDM cosmologies the logarithmic slope of the initial power spectrum ranges between $-3<n<1$, depending on the mass scale (see \\fig~\\ref{fig:pk}). From dwarfs to superclusters the logarithmic slope of the core density profile should range between $0 \\simlt \\alpha \\simlt 2$. As shown in \\fig~\\ref{fig:pk}, the power spectrum in the 1 \\Mpc box is a power law with $n \\approx -2.7$.  According to the aforementioned scaling relationship $\\alpha \\approx 0.4$. This slope is in agreement with what I find in the 1 \\Mpc box.  In the 32 \\Mpc box the power spectrum is not a single power law but the slope is close to $n \\approx -2 $ and therefore $\\alpha \\approx 1$ is expected, consistent with the NFW universal profile and with the result found for this simulation. In cluster and supercluster size halos $\\alpha$ should be larger than in galaxy size halos ($\\alpha \\sim 1.5$).  Indeed, in the 256 \\Mpc simulation I find slopes of the inner profile $\\alpha \\sim 1.4$. But this is not a strong result because the number of particles in each halo is barely sufficient to start seeing a deviation with respect to the NFW profile.  Moreover, the introduction of the parameter $\\alpha$ for the slope of the inner part of the halo profile (which can be easily calculated knowing the mass of the galaxy halo) does not increase the number of adjustable parameters in the profile fitting formula. This happens because the concentration parameter, $c_\\Delta^{NFW}(M_{dm})$, needed to fit the profiles with the NFW formula, is in this case a constant.  The changing slope of the inner profile mimics the variation of the concentration parameter in the NFW profile, shifting the value of the radius where the circular velocity reaches its maximum value (see appendix~\\ref{app}).  \\begin{table*}% \\centering \\caption{Clustering and time steps in the four simulations. $a$ is the scale factor; $N_p$ has two columns: the numbers on the left show the number of particles in the most massive halo, the right contains the mean number of particles for the ten most massive halos.} \\label{tab:one} \\begin{tabular*}{14cm}[]{l|cc||l|cc||l|cc||l|cc} \\multicolumn{3}{c}{1 \\Mpc} \\vline \\vline & \\multicolumn{3}{c}{32 \\Mpc} \\vline \\vline & \\multicolumn{3}{c}{256 \\Mpc} \\vline \\vline & \\multicolumn{3}{c}{1024 \\Mpc} \\\\ \\hline a & steps & $N_p/10^4$ & a & steps & $N_p/10^4$ & a &steps & $N_p/10^4$ & a &  steps & $N_p/10^4$ \\\\ \\hline 0.07 & 1213 & 5.0 2.5 & 0.200 & 1530 & ..  ..  & 0.9 & 2419 & 4.0 3.0 & 2.3 & 1530 & 0.16 0.14\\\\ 0.08 & 1443 & 5.5 3.3 & 0.230 & 1814 & 5.5 4.0 & 1.0 & 2546 & 4.0 3.0 & 4.0 & 1681 & 0.17 0.15 \\\\ 0.085& 1542 & 10 5.0  & 0.250 & 1974 & 4.5 4.0 & 1.2 & 2800 & 5.0 4.0 & 10 & 1809 & 0.17 0.15\\\\ 0.090& 1678 & 8.3 5.3 & 0.270 & 2122 & 13 6.0  & 1.5 & 3078 & 7.5 5.0 & 100 & 1880 & 0.17 0.15\\\\ 0.095& 1804 & 13 6.5  & 0.300 & 2470 & 18 7.5  & 1.7 & 3196 & 7.5 5.5 &1000 & 1910 & .. .. \\\\ \\end{tabular*} \\end{table*} This paper is organised as follows: in \\S~\\ref{sec:sim} I explain the methods for the simulations and in \\S~\\ref{sec:res} I show the results for dwarf galaxies, normal galaxies and clusters of galaxies. In \\S~\\ref{sec:prof} I demonstrate how a density profile with changing inner slope $\\alpha$ can fit all the profiles from dwarfs to clusters with a constant concentration parameter. The discussion and conclusions are presented in \\S~\\ref{sec:con}.  The equations for the circular velocity and integrated mass for the NFW profile and for a halo with arbitrary inner and outer slopes of the density profile are contained in appendix~\\ref{app}. ",
        "conclusions": "\\label{sec:con}  Two main problems faced by CDM cosmologies are related to the properties of small mass galaxies. Namely, (i) the number of visible galactic satellites (dwarf galaxies) in the Local Group is smaller than predicted by N-body simulations \\citep{Moore:99, Klypin:99} and (ii) the flatness of the DM cores in dwarf and LSB galaxies is not reproduced by N-body simulations. The first problem can be solved by including radiative feedback effects \\citep{Chiu:01, RicottiGSb:02} in cosmological simulations. Galaxies with masses $M_{dm} \\simlt 10^8-10^9$ M$_\\odot$ are too small to retain the gas that is photoevaporated by stars or by the ionising background after reionization. Consequently, the luminosity of most dwarf galaxies is predicted to be too low to be detected (\\ie, the mass-to-light ratio should increase for halos with smaller masses) or zero for very small mass halos. Here I have shown results of N-body simulations suggesting that the second problem might not contradict the predictions of CDM cosmologies. But further investigation is needed to understand if the dwarf galaxies showing shallow DM profiles at $z \\sim 10$ (\\eg, just after the time when most of the dwarf galaxies formed) are in a configuration that will survive unchanged until redshift $z=0$. In summary the main findings of this work are: \\begin{itemize} \\item The slope, $\\alpha$ of the inner profile of DM halos is determined by the mass function of the accreting substructure (\\ie, the initial power spectrum of initial perturbations). Dwarf galaxies have on average flatter DM cores, and clusters, steeper cores than galaxies similar to the Milky Way for which the NFW profile is a good fit to the halo density profile.  \\item The logarithmic slope, $\\gamma$, of the outer parts of the halo appears to depend on the acceleration of the universe: when the scale parameter is $a \\simlt 1$, $\\gamma \\approx 3$ as in the NFW profile, but $\\gamma \\approx 4$ at $a >1$ when $\\Omega_\\Lambda \\sim 1$ and the universe is inflating.  \\item A density profile in the form $\\rho_{dm} \\propto X^{-\\alpha}(1+X)^{-\\gamma+\\alpha}$, where $X=c_\\Delta r/r_\\Delta$, $\\alpha$ is a function of the halo mass ($\\alpha \\approx 1.3+[M_{dm,14}^{1/6}-1]/[M_{dm,14}^{1/6}+1]$, where $M_{dm,14}=M_{dm}/[3 \\times 10^{14}$ M$_\\odot$]) and $\\gamma \\sim 3$ (but $\\gamma \\sim 4$ if $\\Omega_\\Lambda \\sim 1$), can fit all the profiles from dwarfs to superclusters with constant concentration parameter $c_\\Delta \\simeq 7$.  This provides a physical explanation for the variation of the concentration parameter in the NFW profile as a function of the halo mass: the changing slopes mimic the variation of the concentration parameter by shifting the radius where the circular velocity reaches its maximum value.  \\item The dependence of $\\alpha$ on the mass of the halo, and therefore on the slope of the power spectrum, $n$, agrees with the theoretical expectation $\\alpha = (9+3n)/(5+n)$, derived assuming that ``undigested'' satellites determine the halo structure in the inner regions \\citep{SubramanianCO:00}. Note that if the DM halo develops a flat core, the chance for accreting satellites to survive tidal disruption is larger \\citep{Dekel:03}. \\end{itemize}  The shape of the profiles for a given mass and redshift present significant statistical variations. This can be attributed to two factors. Firstly, galaxies are not isolated entities and especially at high redshift are severely disturbed by neighbouring satellites or by undergoing major mergers. This produces irregularities in the azimuthally averaged density profile. Secondly, the accretion history of satellites has statistical variation and depends on the local environment. A dwarf galaxy forming at high redshift from a large $\\sigma$ peak of the power spectrum of initial perturbations (the type of galaxy simulated in this work) could have an inner profile different from the same mass galaxy forming at lower redshift from a smaller $\\sigma$ perturbation. Also, the slope of the profile could differ if a galaxy forms in a void or in an overdense region, even for galaxies with similar masses and radii.  I have identified three main reasons to explain why previous studies \\citep[\\eg,][]{Eke:01, Moore:99} found that the inner slope of the DM profiles are not determined by the power spectrum or accretion history (but see also \\cite{SyerW:98, Kravtsov:98, JingS:00, Jing:00} for a different point of view). (i) The profiles of halos with typical masses resolved by the 1 \\Mpc simulation have not been analysed before. (ii) The dynamical state of the the halos at $z=10$ might be different at $z=0$. But there are not published high-resolution simulations for dwarf-size halos ($M_{DM} \\sim 10^8$ M$_\\odot$) at $z=0$. The classic results of \\cite{Cole:96, Navarro:97} are based on the same type of simulations presented here. (iii) Most of the computational efforts have focused on simulating a single galaxy-size or cluster-size halo with mass $M_{dm} \\simgt 10^{12}$ M$_\\odot$ at $z=0$ with higher resolution (about $10^6$ particles per halo) than in this study. Such a large number of particles is required to study the profile in the very inner regions of the halo. But the task of achieving reliable profiles with such a high spatial resolution is sensitive to numerical integration errors and requires careful resolution studies \\citep{Power:02}; perhaps this is a reason for the disagreement between groups, using different codes, on the slope of the inner profile.  In this work I find that the shape of the DM profiles present significant scatter around the mean. It is, therefore, important to analyse a significant statistical sample of halos in order to measure the mean profile.  Works where a single halo is re-simulated with higher resolution are affected by a bias, difficult to control, determined by the criteria for picking the halos to re-simulate. Finally, the fitting method could be important as well for the results. The degeneracy between the slope of the inner profile and the value of the core radius is more easily broken fitting the circular velocities instead of the density profiles. This is because the radius where the circular velocity is maximum is univocally determined by the value of the core radius, $r_s$ (for the NFW profile is $r_{max}=2.16 r_s$). The halo profiles at any redshift and mass, when renormalised matching the values and radii of the maximum circular velocities, have to be identical if their profiles can be all fitted by the NFW profile.  \\cite{Moore:99} find steep profiles in cluster size halos ($\\alpha \\sim 1.5$), in agreement with the results of this work. The results of this work also agree with the classic NFW result for Milky Way size galaxies but not with \\cite{Moore:99} for halos with the same mass. In this work the spatial resolution is not as large as in the aforementioned studies, therefore I cannot study the slope of the DM profile in the very centre of the halos. Nevertheless, I find evidence for disagreement with the NFW predictions even at relatively large radii. Perhaps this new result has been found because it is the first time that a simulation of halos with masses typical of dwarf galaxies and mass resolution $M_p \\sim 10^3$ M$_\\odot$ has been performed. The result for the cluster size halos could also be understood with similar reasoning (\\ie, the cluster masses are larger than in previous studies: $M_{dm} \\sim 10^{15}-10^{16}$ M$_\\odot$).  The steeper core density profile found in clusters may not be as evident as the flatter core found in dwarf galaxies but, combined with the result for normal galaxies, the findings from the three simulations strengthen the case for an inner slope that changes with the mass of the halos.  I conclude this work with two final remarks: \\begin{itemize} \\item[-] According to the scaling relation $\\alpha = (9+3n)/(5+n)$, the inner slope of DM profile of dwarf galaxies is more sensitive to the power spectrum index, $n$, than more massive galaxies: $\\delta\\alpha = 1.5 \\delta n$ if $n \\sim -3$. A tilt $\\delta n \\approx 0.2$ in the initial power spectrum should produce a variation $\\delta \\alpha \\approx 0.3$ of the inner slope of the DM profile.  \\item[-] The best place to look for the effects of the cosmological constant is the outer profile of clusters and superclusters at low redshift. The prediction is for a slope steeper than $\\gamma =3$. Perhaps gravitational lensing studies could be useful in this regard. \\end{itemize}  \\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.  I would like to thank Jerry Ostriker for his valuable comments on the draft and exciting discussions.  \\appendix"
    },
    "0212/astro-ph0212236_arXiv.txt": {
        "abstract": "We present deep near-infrared (NIR) \\vltj, \\vlth, and \\vltk-band ISAAC imaging of the WFPC2 field of the \\hdfslong\\ (\\hdfs). The $2.5\\arcmin \\times 2.5\\arcmin$ high Galactic latitude field was observed with the VLT under the best seeing conditions with integration times amounting to \\jtime\\ hours in \\vltj, \\htime\\ hours in \\vlth, and \\ktime\\ hours in \\vltk. We reach total AB magnitudes for point sources of \\jlimthree, \\hlimthree, and \\klimthree\\ respectively ($3\\sigma$), which make it the deepest ground-based NIR observations to date, and the deepest \\vltk-band data in any field. The effective seeing of the coadded images is $\\approx0\\farcs45$ in $J_s$, $\\approx0\\farcs48$ in $H$, and $\\approx0\\farcs46$ in $K_s$. Using published WFPC2 optical data, we constructed a \\vltk-limited multicolor catalog containing \\catkcounts\\ sources down to $K_{s,AB}^{tot} \\lesssim 26$, of which \\wmincounts\\ have seven-band optical-to-NIR photometry. These data allow us to select normal galaxies from their rest-frame optical properties to high redshift ($z \\lesssim 4$). The observations, data reduction and properties of the final images are discussed, and we address the detection and photometry procedures that were used in making the catalog. In addition, we present deep number counts, color distributions and photometric redshifts of the \\hdfs\\ galaxies. We find that our faint $K_s$-band number counts are flatter than published counts in other deep fields, which might reflect cosmic variations or different analysis techniques. Compared to the HDF-N, we find many galaxies with very red $V-H $ colors at photometric redshifts $1.95 < z_{phot} < 3.5$. These galaxies are bright in $K_s$ with infrared colors redder than $J_s-K_s > 2.3$ (in Johnson magnitudes).  Because they are extremely faint in the observed optical, they would be missed by ultraviolet-optical selection techniques, such as the U-dropout method. ",
        "introduction": "In the past decade, our ability to routinely identify and systematically study distant galaxies has dramatically advanced our knowledge of the high-redshift universe. In particular, the efficient U-dropout technique \\citep{S96a,S96b} has enabled the selection of distant galaxies from optical imaging surveys using simple photometric criteria. Now more than 1000  of these Lyman break galaxies (LBGs) are spectroscopically confirmed at $z \\gtrsim 2$, and have been subject to targeted studies on spatial clustering \\citep{Gi01}, internal kinematics \\citep{Pe98,Pe01}, dust properties \\citep{A00}, and stellar composition \\citep{Sh01,Pa01}. Although LBGs are among the best studied classes of distant galaxies to date, many of their properties like their prior star formation history, stellar population ages, and masses are not well known.  \\par  More importantly, it is unclear if the ultraviolet-optical selection technique alone will give us a fair census of the galaxy population at $z \\sim 3$ as it requires galaxies to have high far-ultraviolet surface brightnesses due to on-going spatially compact and relatively unobscured massive star formation. We know that there exist highly obscured galaxies, detected in sub-mm and radio surveys \\citep{Sm00}, and   optically faint hard X-ray sources \\citep{Co01,B01} at high redshift that would not be selected as LBGs, but their number densities are low compared to LBGs and they might represent rare populations or transient evolutionary phases. In addition, the majority of present-day elliptical and spiral galaxies, when placed at $z \\sim 3$, would not satisfy any of the current selection techniques for high-redshift galaxies. Specifically, they would not be selected as U-dropout galaxies because they are too faint in the rest-frame UV. It is much easier to detect such galaxies in the near-infrared (NIR), where one can access their rest-frame optical light. Furthermore, observations in the near-infrared allow the comparison of galaxies of different epochs at fixed rest-frame wavelengths where long-lived stars may dominate the integrated light. Compared to the rest-frame far-UV, the rest-frame optical light is less sensitive to the effects of dust extinction and on-going star formation, and provides a better tracer of stellar mass.  By selecting galaxies in the near-infrared \\vltk-band, we expect to obtain a more complete census of the galaxies that dominate the stellar mass density in the high-redshift universe, thus tracing the build-up of stellar mass directly.  \\par In this context we initiated the \\fireslong\\ \\citep[\\fires;][]{Fr00}, a large public program carried out at the {\\em Very Large Telescope} (VLT) consisting of very deep NIR imaging of two selected fields. We observed fields with existing deep optical WFPC2 imaging from the {\\em Hubble Space Telescope} (HST): the WPFC2-field of \\hdfslong\\ (\\hdfs), and the field around the $z\\approx0.83$ cluster MS1054-03. The addition of NIR data to the optical photometry is required not only to access the rest-frame optical, but also to determine the redshifts of faint galaxies from their broadband photometry alone. While it may be possible to go to even redder wavelengths from the ground, the gain in terms of effective wavelength leverage is less dramatic compared to the threefold increase going from the $I$ to $K$-band. This is because the \\vltk-band is currently the reddest band where achievable sensitivity and resolution are reasonably comparable to deep space-based optical data. Preliminary results from this program were presented by \\citet[][ hereafter R01]{Ru01}. \\par  Here we present the full NIR data set of the HDF-S, together with a \\vltk-selected multicolor catalog of sources in the \\hdfs\\ with seven-band optical-to-infrared photometry (covering $0.3 - 2.2\\micron$), unique in its image quality and depth. This paper focusses on the observations, data reduction and characteristic properties of the final images. We also describe the source detection and photometric measurement procedures and lay out the contents of the catalog, concluding with the NIR number counts, color distributions of sources, and their photometric redshifts.  The results of the MS1054-03 field will be presented by \\citet{Fo02} and a more detailed explanation of the photometric redshift technique can be found in \\citet{Ru01,Ru02b}. Throughout this paper, all magnitudes are expressed in the AB photometric system \\citep{Ok71} unless explicitly stated otherwise. ",
        "conclusions": "We have presented the results of the FIRES deep NIR imaging of the WFPC2-field of the HDF-S obtained with ISAAC at the VLT: the deepest ground-based NIR data available, and the deepest \\vltk-band of any field. We constructed a \\vltk-selected multicolor catalog of galaxies, consisting of 833 objects with $K_{s,AB} \\lesssim 26$ and photometry in seven-bands from $0.3$ to $2.2\\mu$ for \\wmincounts\\ of them. These data are available electronically together with photometric redshifts for \\nphotz\\ galaxies. Our unique combination of deep optical space-based data from the HST together with deep ground-based NIR data from the VLT allows us to sample light redder than the rest-frame V-band in galaxies with $z\\lesssim3$ and to select galaxies from their rest-frame optical properties, obtaining a more complete census of the stellar mass in the high-redshift universe. We summarize our main findings below: \\par  \\begin{itemize}  \\item The $K_s$-band galaxy counts in HDF-S turn over at the faintest magnitudes and flatten from $\\alpha \\approx 0.25$ at total AB magnitudes $22 \\lesssim K_s \\lesssim 24$ to $\\alpha \\approx 0.15$ at $24 \\lesssim K_s \\lesssim 26$; this is flatter than counts in previously published deep NIR surveys, where the FIRES HDF-S field is largest and deepest amongst these surveys. The nature of the scatter in the faint-end counts is yet unclear but field-to-field varations as well as different analysis techniques likely play a role.  \\item The HDF-S contains 7 sources redder than $(V_{606} - H)_{AB} \\gtrsim 3$ and brighter than total magnitude $H_{AB} \\lesssim 25.5$ at photometric redshifts $1.95 < z_{phot} < 3.5$, while such galaxies were virtually absent in the HDF-N. They are much redder than regular U-dropout galaxies in the same field and are candidates for relatively massive, evolved systems at high redshift. The difference with the HDF-N might just reflect field-to-field variance, calling for more observations to similar limits with full optical-to-infrared coverage. Results from the second and larger FIRES field centered on MS1054-03 \\citep{Fo02} should provide more insight into this issue.  \\item We find substantial numbers of red galaxies with $(J_s - K_s)_J > 2.3$ that have photometric redshifts $z_{phot} > 2$. These galaxies would be missed by ultraviolet-optical color selection techniques such as the U-dropout method because most of them are barely detectable even in the deepest optical images. The surface densities of these sources in our field keeps rising down to the detection limit in $K_s$, in contrast to the number counts of EROs which peak at $K_{s,AB}\\sim 23$ and then drop or flatten at fainter magnitudes.   \\end{itemize} \\par  The results of the HDF-S presented in this paper demonstrate the necessity of extending optical observations to near-IR wavelengths for a more complete census of the early universe. Our deep $K_s$-band data prove invaluable for they probe well into the rest-frame optical at $2 < z < 4$, where long-lived stars may dominate the light of galaxies. We are pursuing follow-up programs to obtain more spectroscopic redshifts needed to confirm the above results. Updates on the FIRES programme and access to the reduced images and catalogues can be found at our website \\fireswww."
    },
    "0212/astro-ph0212326_arXiv.txt": {
        "abstract": "The {\\it Chandra} discovery of bright X-ray emission from kpc-scale jets provides us unprecedented insights into the physical state of the plasma in the flow. In particular it is possible to get good constraints on the power and pressure in bright knots.  For a group of selected sources with blazar-type cores it is also possible to constrain the physical quantities of the jet at sub-pc scale. We discuss how these results can help us to connect the properties of the jet at different scales. ",
        "introduction": "\\label{intro} The study of extragalactic jets has been recently renewed by the discovery, made by {\\it Chandra}, of numerous jets bright in the X-rays at scales $>$ kpc, both in FRI and FRII sources (Chartas et al. 2000, Sambruna et al. 2002). While in low-power FRI jets the dominating emission mechanism responsible for X-rays is thought to be synchrotron from very-high energy electrons (e.g. Worrall et al. 2001), in the case of the most powerful sources it is clear that X-rays constitute a second spectral component, likely due to the IC scattering of the CMB photons (Tavecchio et al. 2000a, Celotti et al. 2001), implying that X-ray bright jets are still relativistic (with bulk Lorentz factors $5-10$) at $\\sim 100$ kpc.\\\\ At the opposite side of scalelengths, studies of the innermost portion ($d\\sim 0.1$ pc) of jets in AGNs are rather well developed.  Our knowledge of these regions of jets is mainly based on blazars, whose Spectral Energy Distribution (extending from radio to $\\gamma$-rays) is dominated by the relativistic amplified nonthermal continuum produced in a jet closely aligned to the line of sight (e.g. Urry \\& Padovani 1995). The double humped SED of these sources is well described by synchrotron-IC models (e.g., Ghisellini et al. 1998; Sikora \\& Madejski 2001).\\\\ The possibility to constrain the physical state of the plasma in the jet both at sub-pc and kpc scale offers us the interesting opportunity to shed some light on the evolution of the jet from very small scales, close to the central engine, to the outer regions, where the jet is starting to significantly decelerate. In this work we present a first step, reporting the results obtained for a small group of sources for which good data for both regions are available. These sources are PKS~1510-089, 1641+399, PKS~0521-365 and 3C371. The emission models from the blazar region and the parameters for the first three sources are discussed in Tavecchio et al. (2000b, 2002), while the model for 3C371 is discussed in Tavecchio et al. (in prep). Multiwavelength ({\\it Chandra}, {\\it HST} and radio) data for extended jets are discussed in Gambill et al. (in prep 1510-089; 1641+399), Pesce et al. (2001, 3C371) and Birkinshaw et al. (2002, 0521-365). Clearly the X-ray emission for the first two sources is well reproduced by the IC/CMB model, while in the other two jets the dominant contribution is due to synchrotron emission. The detailed modelling of the data will be reported in Tavecchio et al. (in prep). In the following we limit the discussion to the result of the comparison of the most relevant quantities of the jet as inferred in the two regions. ",
        "conclusions": "The results discussed above seems to point toward a simple scenario for the dynamical evolution of the jet from inner to the outer regions. Clearly the pressure in the blazar jet is larger than any possible confining gas. This means that, after the dissipation region, the jet expands freely until the internal pressure (gas and/or magnetic) reaches the pressure of the (supposed) hot gas of the galactic halo. At this point a reconfinement shock forms (the jet is highly supersonic), particle are accelerated, and the jet starts to decelerate (Sanders 1983, Komissarov 1994).  This scenario seems to account for the typical distance at which knots form and also provide a reliable mechanism for the production of relativistic particles. A even more interesting result is that the expected pressure decay along the jet seems to be consistent with the simple profiles predicted for a free, conical expanding jet (Blandford \\& Koenigl 1979).  A similar scenario has been recently proposed to explain the dynamics of the jet in the FRI radio galaxy 3C31 by Laing \\& Bridle (2002). The results discussed here are based on a very small group of sources: future studies along these lines could further support this scenario, providing important clues on the dynamics of relativistic jets."
    },
    "0212/astro-ph0212110_arXiv.txt": {
        "abstract": "It is shown that the number of members, $N_A$, and the linear radius of the Abell clusters at $z<0.18$ increase with redshift. The increase is caused by observational selection: the number of rich and large clusters is small in the nearby small volume of space, while poor and small clusters are not detected at high distances. Besides, the numbers of clusters with different ellipticities are about the same at all distances, though in the case of cosmological evolution the number of the more elliptical clusters increases with decrease of distance. It means that clusters do not evolve dynamically. The dependences of the X-ray emission detection rate and the temperature of the X-ray emission on redshift are consequences of the dependence of $N_A$ on redshift. The radial velocity dispersion of the less elliptical clusters is higher than that of the elongated ones. This shows that galaxies in a cluster move preferentially along the large axis of the cluster. The space density is not changing in the considered range of $z$. Movement of galaxies along the cluster elongation may be explained by rotation along the large axis of a cluster. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/astro-ph0212260_arXiv.txt": {
        "abstract": "Statistical parameters of the ISM driven by thermal energy injections from supernova explosions have been obtained from 3D, nonlinear, magnetohydrodynamic, shearing-box simulations for spiral arm and interarm regions. The density scale height obtained for the interarm regions is 50\\% larger than within the spiral arms because of the higher gas temperature. The filling factor of the hot gas is also significantly larger between the arms and depends sensitively on magnetic field strength. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/astro-ph0212056_arXiv.txt": {
        "abstract": " ",
        "introduction": "The Compton Gamma Ray Observatory has detected seven $\\gamma$-ray pulsars. The observed light curves and energy spectra have been used to discriminate possible radiation models. Moreover, the next-generation $\\gamma$-ray space telescopes and the ground based Cherenkov telescopes will further constrain to the models.  The pulsed $\\gamma$-rays imply that the electrons and positrons are accelerated up to about 10 TeV in the magnetosphere. For the outer-gap model (Chen, Ho \\& Ruderman, 1986), Hirotani \\& Shibata (1999, HS) solved the accelerating electric field with curvature radiation and pair creation processes self-consistently. They showed that the electric field along the magnetic field does not extend to the light cylinder, and calculated curvature spectra can explain the EGRET observations in the GeV bands. However, the gap emission could not explain the observations in the MeV band. Hence, to improve the HS model, we take account of the radiation from the outside of the gap and compare the corrected spectrum with the observations. ",
        "conclusions": "In summary, we obtained the spectrum in good agreement with the EGRET observations for Vela by inclusion of the curvature radiation from the outside of the gap. We found that the observed peak energies for Vela and B1706-44 may imply that the almost the same currents in units of GJ value are running through the gap for both pulsars.  From Fig.2 (a), we recognize that the calculated spectrum is inconsistent with the COMPTEL observations. If we try to explain this observations, we need the very small curvature radius as compared with the dipole, which is unlikely. Therefore, this MeV emission will be obtained by inclusion of the synchrotron emission by pairs.  In \\S\\S4.2, we showed that we need very large $A_{\\perp}$ to explain the observations of B1706-44. This may be due to the small $j_{tot}$, about 20\\% of the GJ current. However, if we adopt the nearly ($j_{tot}\\sim1$) or super ($j_{tot}>1$) GJ current with this model, the gap width will be quenched.  Quite recently, Hirotani et.al. (2002) showed that the value of the $\\Gamma_{sat}$ given by eq.(\\ref{lorent}) and also $\\gamma$-ray fluxes calculated with this $\\Gamma_{sat}$ are overestimated, because particles in the gap do not immediately saturate. But, the general features of the gap model and the result that we must take account of the radiation from the outside of the gap to explain the observations above 100MeV are not altered.   $Acknowledgments$\\quad We thank Drs. Tanimori and Kushida for useful discussion for B1706-44 and the CANGAROO group for holding the symposium."
    },
    "0212/astro-ph0212506_arXiv.txt": {
        "abstract": "% The interstellar abundances of D I, N I, and O I in the local ISM are studied using high-resolution spectra of four hot white dwarfs. The spectra of GD 246, WD 2331$-$475, HZ 21, and Lan 23 were obtained with the {\\it Far Ultraviolet Spectroscopic Explorer} (\\fuse) in the wavelength range 905~--~1187~\\AA. The line of sight to GD 246 probes the Local Interstellar Cloud and at least one other H I cloud inside the Local Bubble, which contains most of the gas seen along this line of sight. The column densities of H I, C II*, S II, and Si II are measured using archival {\\it Hubble Space Telescope} STIS echelle-mode observations. The H I column density is determined by fitting the strong damping wings of interstellar \\lya~using a model atmosphere to account for the stellar continuum. The sightline-averaged ratios for GD 246 are: D~I/H~I~=~(1.51~$\\pm~^{0.39}_{0.33})\\EE{-5}$, O~I/H~I~=~(3.63~$\\pm~^{0.77}_{0.67})\\EE{-4}$, and D~I/O~I~=~(4.17~$\\pm~^{1.20}_{1.00})\\EE{-2}$  (uncertainties are 2$\\sigma$). This line of sight provides the fourth reliable \\fuse~measurement of the Local Bubble D/H ratio. For the WD 2331$-$475 line of sight we find sightline-averaged ratios: D~I/O~I~=~(5.13~$\\pm~^{2.20}_{1.69})\\EE{-2}$ and D I/N I~=~(4.57~$\\pm~^{1.88}_{1.45})\\EE{-1}$. Toward HZ 21 the sightline-averaged ratios are: D~I/O~I~=~(4.57~$\\pm~^{2.22}_{1.63})\\EE{-2}$ and D I/N I~=~(4.27~$\\pm~^{1.96}_{1.44})\\EE{-1}$. In the higher column density sightline to Lan 23 the sightline-averaged ratios are: D~I/O~I~=~(3.24~$\\pm~^{3.27}_{2.06})\\EE{-2}$ and D~I/N~I~=~(3.16~$\\pm~^{1.56}_{1.23})\\EE{-1}$. Molecular hydrogen, corresponding to rotational levels J $\\le$ 3, is clearly seen along this line of sight. No reliable H I measurements are available for WD 2331$-$475, HZ 21, or Lan 23. We combine the different abundance ratios computed here with previous published values to produce revised \\fuse~abundance ratios for D I/H I, O I/H I, N I/H I, D I/N I, D I/O I, and O I/N I. ",
        "introduction": "%  The present day abundance ratio of deuterium to hydrogen places important constraints on Big Bang nucleosynthesis (BBN) and the chemical evolution of galaxies. Since it is believed that deuterium is only produced in appreciable amounts in primordial BBN and destroyed in stellar interiors, the measurement of D I/H I in the interstellar medium (ISM) places a lower limit on the primordial abundance of deuterium. By comparing the ISM abundance of deuterium to its abundance in high-redshift intergalactic gas we should be able to understand better the effects of astration and chemical evolution of galaxies. Measurements of the D/H ratio in intervening clouds of gas seen toward distant quasars have yielded a range of values D/H = (1.65 -- 4.0)$\\EE{-5}$ \\citep[and references therein]{2001ApJ...552..718O,2001ApJ...560...41P,2002ApJ...565..696L}. Measurements of D/H in the local ISM have been made with~\\copernicus~\\citep[e.g.][]{1973ApJ...186L..95R}, \\hst, \\citep[e.g.][]{1995ApJ...451..335L},~\\imaps,~\\citep{1999ApJ...520..182J,2000ApJ...545..277S}, and more recently~\\fuse~\\citep[and references therein]{2002ApJS..140....3M}. A nearly constant ratio of D/H = $(1.5~\\pm~0.1)\\EE{-5}$ (1$\\sigma$~on the mean) has been obtained in the Local Interstellar Cloud (LIC) by \\citet{1998SSRv...84..285L}; recent measurements inside the Local Bubble \\citep{2002ApJS..140....3M} appear to be consistent with a single value for D/H in the Local Bubble. Other measurements \\citep{1979ApJ...229..923L,1983ApJ...264..172Y,1999A&A...350..643H,1999ApJ...520..182J,2000ApJ...545..277S} suggest variations of the interstellar D/H ratio beyond the Local Bubble, at the distance of a few hundred parsecs.  Until \\fuse~was launched (1999 June 24) only the UV spectrographs onboard \\hst~could be used for systematic studies of the deuterium abundance in the ISM. Since only the \\lya~transitions of H I and D I are observable in the \\hst~bandpass, such studies were restricted to low column density sightlines so that D I was neither completely obscured by the adjacent H I \\lya~nor heavily saturated. With \\fuse, we now have access to the complete Lyman series of deuterium (except \\lya), which allows the study of a larger range of environments and provides tighter constraints on the abundance of D I compared to studies using only \\lya. In addition, the numerous lines of N I, O I, and Fe II in the \\fuse~bandpass can be used to trace the metallicity and dust content of the absorbing gas.   In this work we present the first measurements of deuterium absorption toward the white dwarfs GD 246, WD 2331$-$475, HZ 21, and Lan 23. Table \\ref{stellar_prop} summarizes the stellar parameters. This study adds four more sightlines to the seven sightline studies that constitute the first set of \\fuse~deuterium measurements \\citep[see overview paper by][and references therein]{2002ApJS..140....3M}. This paper is organized as follows. The observations and data processing are presented in \\S 2; the analysis methodology is described in \\S 3. The GD 246 line of sight analysis is presented in \\S 4. The WD 2331$-$475, HZ 21, and Lan 23 lines of sight are analyzed in \\S 5, and systematic effects are discussed in \\S 6. We discuss the results in \\S 7. ",
        "conclusions": "In section 7.1 we discuss the line of sight to the four stars in the context of the Local Bubble. In \\S~7.2 we compare the total column density ratios of D/H, O/H, and N/H for GD 246, WD 2331$-$475, HZ 21, and Lan 23 to values found for other lines of sight in the local ISM and discuss the implications of using O and N as tracers of H for these lines of sight. In \\S~7.3 we combine the ratios derived here with previously published \\fuse~ratios, deriving revised \\fuse~mean ratios. In \\S~7.4 we discuss the mean D/O ratio inside the Local Bubble. \\S~7.5 presents a more detailed discussion of the GD 246 line of sight. Because we have information on the velocity structure along this line of sight we can separate the LIC and Local Bubble contributions to the total column densities and draw more detailed conclusions about the distribution of material along the line of sight.  \\subsection{Local Bubble}  \\citet{1999A&A...346..785S} have used absorption line studies of Na I to approximately map the contours of the Local Bubble. They found a cavity with a radius between 65 and 250 pc (depending on the direction) that is delineated by a sharp gradient in the neutral gas column density with increasing radius, a dense neutral gas ``wall''. \\citet{1999A&A...346..785S} quote an equivalent width of 20~m\\AA~of Na I $\\lambda$5891.59 as corresponding to log N(H I)~$\\sim$~19.3, which is used as a rough estimate of a plausible Local Bubble boundary. According to their map of the Local Bubble, the GD 246 line of sight does not penetrate the wall, as it lies between the 5 m\\AA~and 20 m\\AA~Na I $\\lambda$5889.95 contours. Although the uncertainty in the distance to GD 246 (and for the other three white dwarfs, for that matter) makes this result ambiguous, log N(H I)~=~19.11~$\\pm$~0.05 implying that GD 246 is likely close to the wall of the Local Bubble, but does not penetrate the higher density regions beyond. The lines of sight to WD 2331$-$475 and HZ 21 do not penetrate the wall either; WD 2331$-$475 lies inside the 20 m\\AA~contour and HZ 21 lies between the 5 m\\AA~and 10 m\\AA~contours. In addition, as shown in the next section, the measured O I column density toward both WD 2331$-$475 and HZ 21 implies an H I column density similar to that towards GD 246. The Lan 23 line of sight is located between the 20 m\\AA~and 50 m\\AA~Na I $\\lambda$5889.95 contours, and has a high H I column density; it likely penetrates the wall of the Local Bubble.   \\subsection{O/H, D/H, and N/H Ratios}  Table \\ref{ratios} summarizes the column density ratios for several species along the four lines of sight studied in this work. For GD 246 we find O I/ H I~=~(3.63~$\\pm~^{0.77}_{0.67}$)$\\EE{-4}$. This ratio is in agreement with the O I/H I ratio derived by \\citet{1998ApJ...493..222M}, of O I/H I~=~(3.43~$\\pm$~0.30)$\\EE{-4}$~(using the $f$-value updated to $f$~=~1.16$\\EE{-6}$~from \\citet{1999ApJS..124..465W}.  Reliable H I measurements are not available for WD 2331$-$475, HZ 21, or Lan 23. However, the constancy of the O/H ratio discussed above enables us to derive an approximate column density for H I using O I. Using the \\citet{1998ApJ...493..222M} interstellar O I abundance  and the O I column density along each line of sight, we estimate log N(H I)~=~18.94 for WD 2331$-$475, log N(H I)~=~19.20 for HZ 21, and log N(H I)~=~20.16 for Lan 23.  Using the N(H I) obtained this way and our measured N(D I) we obtain D I/H I~$\\sim$~1.8$\\EE{-5}$, 1.6$\\EE{-5}$, and 1.2$\\EE{-5}$, for WD 2331$-$475, HZ 21, and Lan 23, respectively. The uncertainties in the O I columns used to derive the H I column densities are large enough that the differences between the derived D/H ratios for these stars and the \\citet{2002ApJS..140....3M} weighted mean of D/H~=~(1.52~$\\pm$~0.15)$\\EE{-5}$ are not significant. The D/H ratio for GD 246 is discussed below in \\S~7.5.    Using the \\citet{1997ApJ...490L.103M} interstellar abundance of nitrogen, N/H~=~(7.5~$\\pm$~0.8)$\\EE{-5}$, and our measured N I column density leads to estimated values of log N(H I)~=~18.87 for GD 246, log N(H I)~=~18.65 for WD 2331$-$475, log N(H I)~=~18.89 for HZ 21, and log N(H I)~=~19.82 for Lan 23. The H I column density derived using the N I column density is in all four cases only $\\sim$50$\\%$~of the H I column density determined using the O I abundance, implying that the N I column densities are low. It is possible that the N along these lines of sight is $\\sim$50$\\%$~ionized. This effect has been seen for other sightlines as well \\citep[see][]{2000ApJ...538L..81J,2002ApJS..140...81L,2002ApJS..140....3M}. For GD 246 we measure N I/H I~=~(4.37~$\\pm~^{0.84}_{0.74}$)$\\EE{-5}$ which is a factor of 1.7 times smaller than the \\citet{1997ApJ...490L.103M} value, but is in close agreement with the weighted mean of the \\fuse~sightlines N I/H I~=~(4.24~$\\pm$~0.62)$\\EE{-5}$ \\citep[see][]{2002ApJS..140....3M}.  The ionization fraction of N I along a line of sight can be estimated directly by measuring the N II and N III column densities as both N II and N III have transitions in the \\fuse~bandpass. However, for our sightlines the N II transitions available in the \\fuse~bandpass are saturated, and N III cannot be determined since it is blended with interstellar Si II and sometimes with photospheric N III. \\citet{2002ApJS..140..103H,2002ApJS..140...19K,2002ApJS..140...91W} have been able to measure the column density of N II (and sometimes N III) along three different lines of sight and found that in all cases it was larger than the column density of N I, providing a better agreement between (N/H)$_{total}$ within 100 pc and N I/H$_{total}$ for larger path lengths.  \\subsection{Revised \\fuse~Ratios}  In Table \\ref{comparison} we combine the ratios derived here with the \\citet{2002ApJS..140....3M} ratios for D I/H I, O I/H I, N I/H I, D I/N I, D I/O I, and O I/N I, and present the \\citet{1997ApJ...490L.103M,1998ApJ...493..222M} N/H and O/H values for comparison. We present also the $\\chi^{2}_{\\nu}$ test for a single mean value. The values quoted in Table \\ref{comparison} are the weighted means and uncertainty in the means, where we use the largest of the lower and upper error bars of each individual ratio to compute the weighted mean. {\\it We use 1$\\sigma$~error bars in this subsection} in order to make it easier to compare the different lines of sight among themselves. For Lan 23 values only the D I/N I value is used, due to the large error bars for the other ratios associated with this line of sight . By combining our ratios with the \\citet{2002ApJS..140....3M} values, the weighted means increase by 0 to 5\\%.  The mean D I/O I ratio is (4.06~$\\pm$~0.17)$\\EE{-2}$ with $\\chi^{2}_{\\nu}$~=~1.9 ($\\nu$~=~9). The mean D I/N I ratio now computed with ten lines of sight is (3.41~$\\pm$~0.15)$\\EE{-1}$ with $\\chi^{2}_{\\nu}$~=~2.2 ($\\nu$~=~9).  \\subsection{D/O Ratio Inside the Local Bubble}  \\citet{2001IAP} and \\citet{2002ApJS..140....3M} have found a nearly constant D/O ratio inside the Local Bubble, indicating that D/O traces the D/H ratio.  The weighted mean of the five \\fuse~sightlines within the Local Bubble \\citep{2002ApJS..140....3M} is (3.76~$\\pm$~0.20)$\\EE{-2}$~(1$\\sigma$ in the mean). When we combine the D/O ratios derived here (Lan 23 value not included, as it is likely outside the Local Bubble) with the Local Bubble \\fuse~results from Moos et al. (2002), we derive a weighted mean and uncertainty in the mean of D I/O I~=~(3.87~$\\pm$~0.18)$\\EE{-2}$~(1$\\sigma$ in the mean) with $\\chi^{2}_{\\nu}$~=~0.96 ($\\nu$~=~7), reinforcing the Moos et al. (2002) conclusion that the variability of the data is consistent with the uncertainties, i. e., the D I/O I ratio is constant within the Local Bubble. The D I/O I ratio inside the Local Bubble will be discussed in more detail by H\\'{e}brard et al. (2002, in preparation).   \\subsection{GD 246}  Absorption by at least two different components, separated by $\\sim$9 km s$^{-1}$, is clearly seen, along the GD 246 sightline, in the high resolution STIS echelle data of N I, S II, and Si II (see Figure \\ref{stisdata}). The stronger component (component 1), seen at $-$9.76 km s$^{-1}$, contains 86$\\%$, 82$\\%$, and 79$\\%$ of the total column densities of N I, S II, and Si II, respectively, along this line of sight. The velocity of component 2, $v_{2}$~=~$-$0.63 km s$^{-1}$, is consistent with the expected velocity of the Local Interstellar Cloud (LIC) along this line of sight \\citep[+2.46~$\\pm$~1.84 km s$^{-1}$;][]{1995A&A...304..461L} when the zero point velocity uncertainty in STIS E140H data is taken into account ($\\Delta$$v$~$\\sim$~0.7 -- 1.5 km s$^{-1}$). The discrepancy between the expected LIC velocity and the measured velocity along this sightline might be an indication that the LIC velocity field is more complicated than a single vector. A similar effect has been seen toward the Hyades Cluster by \\citet{2001ApJ...551..413R}, who found $v_{\\rm{predicted}}$(LIC) $-$ $v_{\\rm{observed}}$(LIC) = 2.9 $\\pm$ 0.7 km s$^{-1}$, similar to our value $v_{\\rm{predicted}}$(LIC) $-$ $v_{\\rm{observed}}$(LIC) $\\sim$ 3.1 km s$^{-1}$. The eighteen Hyades stars studied by these authors are located in the sky not too far frm GD 246; for four of them (HD 27561, HD 27848, HD 28237, and HD 29225; $l~\\sim$ 178 -- 184, $b~\\sim$ $-$20 -- $-$25) they also find $v_{\\rm{predicted}}$(LIC) $-$ $v_{\\rm{observed}}$(LIC) $\\sim$ 3.0 km s$^{-1}$.  Sulfur should not be depleted onto dust grains, making it possible to estimate how neutral hydrogen would be distributed between the two components using S II as a proxy for H I. We derive log~N(H I)$_{2}$~=~18.39 using the S II column density for the LIC component (component 2), log~N(S II)$_{2}$~=~13.59. We derive a distance of 8 pc to the edge of the LIC along this direction \\citep[assuming a LIC neutral hydrogen column density of $n_{H I}$~=~0.10 cm$^{-3}$;][]{2000ApJ...528..756L}. Both the H I column density and distance to the edge of the LIC derived here are greater than the values derived from the model of Redfield (2002, private communication), for this direction, of log N(H I)~=~17.93 and d~=~2.75 pc. Our values are closer to the maximum values derived for the LIC by \\citet{2000ApJ...534..825R} of 18.32 and 6.8 pc, respectively. The discrepancy between our values and the LIC model values can be due either to the preliminary nature of their model or to the presence of an unresolved cloud close to the LIC velocity. While H I may not follow exactly the same distribution as S II it is still clear that most of the neutral gas along this line of sight is outside the LIC but within the Local Bubble boundaries.  We find D~I/H~I~=~(1.51~$\\pm~^{0.39}_{0.33}$)$\\EE{-5}$ for this line of sight (2$\\sigma$~error bars). So far, only three lines of sight studied by \\fuse~(HZ 43A, G191-B2B, and WD 1634$-$573) probing the Local Bubble have accurate D/H measurements \\citep{2002ApJS..140...19K,2002ApJS..140...67L,2002ApJS..140...91W}. This is the fourth line of sight, with an accurate D/H measurement that probes the region mentioned above.  The D I/H I ratio for all the gas in this sightline is (1.51~$\\pm~^{0.39}_{0.33}$)$\\EE{-5}$. \\citet{1998SSRv...84..285L} reports a mean value of (1.50~$\\pm$~0.10)$\\EE{-5}$ (1$\\sigma$ in the mean) for 12 LIC sightlines. Because most of the gas along the sightline to GD 246 resides outside the LIC, we conclude that the D I/H I ratio for the gas outside, is also $\\sim$1.5$\\EE{-5}$, similar to the LIC value.   The electron density in component 1, which probes material outside the LIC, can be computed under the assumption that electron collisions populate the upper fine structure level ($J$~=~3/2) of the ground state of C II leading to C II* \\citep[see][]{1993ApJ...409..299S}. Because all the C II lines in the \\fuse~and STIS data are saturated, S II is used as a proxy for C II. This is a reasonable assumption; C II and S II have similar ionization potentials (24.38 eV for C II and 23.33 eV for S II) and should therefore trace the same type of gas. The collisional equilibrium equation is \\begin{equation} A_{21}n(C^{+*})=\\gamma_{12}n_{e}n(C^{+}) \\end{equation} where $\\gamma_{12}$ is the rate coefficient for excitation and A$_{21}$ is the spontaneous downward transition probability. For excitation of ions by electrons \\citep{1978ppim.book.....S}: \\begin{equation} \\gamma_{12}=\\frac{8.63\\EE{-6}}{g_{1}T^{0.5}}\\Omega_{12}~exp(-\\frac{E_{12}}{kT})~cm^3~s^{-1} \\end{equation} $E_{12}$~=~7.9$\\EE{-3}$~eV, and we adopt $A_{21}$~=~2.29 \\citep{1983IAUS..103..143M} and $\\Omega_{12}$~=~2.90 \\citep{1989agna.book.....O}. Assuming that conditions are uniform within the cloud we can replace particle density ratios by column density ratios and assuming that for warm low density clouds the depletions of S and C are 0 and $-$0.3 dex respectively, we find \\begin{equation} \\frac{N(C^{+*})}{N(S II)}=\\frac{55n_{e}}{T^{0.5}}~. \\end{equation} Using the S II column density we measured for component 1, log N(S II)$_{1}$~=~14.25~$\\pm$~0.02, and the C II* column density, log N(C II*)~=~13.05~$\\pm$~0.04, we find n$_{e}$~=~(0.1~$\\pm$~0.01)(T/8000)$^{0.5}$ cm$^{-3}$ (all error bars are 2$\\sigma$). For T~=~4000--8000 K this yields an electron density, $n_{e}$~=~0.06--0.11 cm$^{-3}$ for the material beyond the LIC,  which is similar to values found in the LIC by \\citet{2000ApJ...528..756L}.  Thus, the D/H ratio and the electron density obtained in this study indicate that the properties of the cloud(s) beyond the LIC along the GD 246 line of sight are likely similar to those of the LIC.  The presence of ionized gas in this cloud will not affect the D/H ratio in this region, as D I and H I, with a similar ionization potential, are ionized at the same rate.   Table \\ref{depletions} compares abundances toward GD 246 with solar photosphere abundances. We find there is an indication that N, O, Si, P, Ar, and Fe are depleted from typical solar abundances, but the uncertainties preclude a definitive statement. \\citet{2002ApJS..140....3M} and \\citet{2002SSRLinsky} have noted the evidence for 20 -- 25\\% depletion of O I by dust in the near ISM. For all species listed above part of this depletion might be explained by not taking into account higher ionization states when calculating their gas-phase abundances. N and Si are the only species for which we could measure some of this contribution; however we have only a lower limit on N(N II) and Si III is blended with stellar Si III.  \\subsection{Summary} We have obtained column densities of H I, D I, C II*, N I, O I, Si II, P II, S II, Ar I, and Fe II for the line of sight to the white dwarf GD 246. This is the fourth line of sight with an accurate D/H measurement with \\fuse~that probes the region outside the LIC but inside the Local Bubble. Based on a constant D/H ratio of 1.5$\\EE{-5}$ inside the LIC we find a similar ratio outside the LIC. In addition, our sightline-averaged D I/H I measurement for GD 246 is in agreement with the weighted mean of D I/H I reported by \\citet{2002ApJS..140....3M} of D I/H I~=~(1.52~$\\pm$~0.08)$\\EE{-5}$, for the local ISM. For the WD 2331$-$475, HZ 21, and Lan 23 sightlines the column densities of D I, N I, O I, and other species are also measured and ratios are computed. We have used results from all four sightlines to compute revised \\fuse~ratios for the local ISM."
    },
    "0212/astro-ph0212395_arXiv.txt": {
        "abstract": "\\lya absorber correlations contain information about the underlying density distribution associated with a particular class of absorbers.  As such, they provide an opportunity to independently measure the ``bias\" of the Lyman alpha forest, i.e. the relationship between \\ion{H}{1} column density and underlying dark matter density.  In these proceedings we use hydrodynamic simulations to investigate whether the evolution of this bias is measurable from observable correlations.  Unfortunately, the increasingly complex physics in the IGM at $z\\la 1$ makes a direct measurement of the bias difficult.  Nevertheless, current simulations do make predictions for \\ion{H}{1} absorber correlations that are in broad agreement with observations at both high and low redshift, thus reinforcing the bias evolution predictions given by these models. ",
        "introduction": "\\begin{figure} \\plottwo{fig.dencorr.ps}{fig.cencorr.H1216.ps} \\caption{-- {\\it Left panel:} Line-of-sight density correlations at $z=0,1,2$ for pixels in the range of $-0.5<\\log{\\rho/\\bar{\\rho}}<0.5$, $0.5<\\log{\\rho/\\bar{\\rho}}<1.5$, and $\\log{\\rho/\\bar{\\rho}}>1.5$. {\\it Right panel:} Line-of-sight correlations of pixels with \\ion{H}{1} optical depth $\\tau>0.125,0.5,2$. \\label{fig: dencorr}} \\end{figure}  In bottom-up hierarchical structure formation models, larger overdensities are more strongly correlated.  In Figure~\\ref{fig: dencorr} (left panel) we show how this relationship is manifested along one-dimensional redshift-space lines of sight through our simulation.  We compute the excess number of pairs of pixels having densities (normalized to the cosmic mean) in the ranges typical of the \\lya forest, relative to a randomly distributed set of such pixels.  Despite smearing by peculiar velocities, higher overdensities are still more strongly correlated, though in velocity-space the correlation length is only $\\sim 100$~km/s. Note that the correlation {\\it length} does not increase with density, contrary to what one finds in large-scale structure studies where e.g. the correlation length of clusters is higher than that of galaxies. Also, it is evident that the density correlation doesn't change from $z=2\\rightarrow 0$, indicating stable clustering of \\lya forest structures.  At high redshift, the density is tightly correlated with the optical depth of \\lya absorption.  Hence we would expect that \\lya optical depths would reflect correlations in the density.  The evolution of such correlations would then, in principle, reflect the evolution of bias for a particular optical depth, since the density correlations are not evolving. The right panel of Figure~\\ref{fig: dencorr} shows the correlations for pixels with optical depths $\\tau>0.125,0.5,2$.  According to Cen et al. (1998), this measure most accurately reflects the underlying matter correlation, at least at $z=3$.  Note that here we are using the noise-added spectra, with the optical depth computed by inverting the flux, so the $\\tau$ limits are actually flux limits of $F<0.88,0.61,0.14$, respectively.  Figure~\\ref{fig: dencorr} (right panel) shows that the flux correlations indeed show similar trends as density correlations at $z=2$ and 1, but by $z=0$ the optical depth limits do not clearly delineate density cuts. Furthermore, even from $z=2\\rightarrow 1$ the correlation strength (for $\\Delta v<100$~km/s) has evolved very little, despite a comparatively large evolution in the photoionizing background that governs the bias (see e.g. Figure~7 in Dav\\'e \\& Tripp 2001).  At $z=0$, there is even a significant {\\it anti-correlation} for $\\tau>2$ pixels, until $\\Delta v<100$~km/s.  Thus, disappointingly, it appears that the pixel flux (or optical depth) correlations do not straightforwardly trace the evolution of the \\lya forest bias. ",
        "conclusions": ""
    },
    "0212/astro-ph0212348_arXiv.txt": {
        "abstract": "{A brief review about the chemical composition of cosmic rays in the energy range $10^{15} \\la E \\la 10^{20}$~eV is given. While there is convincing evidence for an increasingly heavier composition above the knee, no clear picture has emerged at the highest energies, yet.  We discuss implications about the origin of cosmic rays and emphasize systematic differences related to data analysis techniques and to the limited understanding of the air shower development.} ",
        "introduction": "When being asked to formulate the most crucial, yet unanswered questions about {\\sc cr}s, one likely would end up with a list as such: \\begin{enumerate} \\item Where and what are the sources of high energy {\\sc cr}s? \\item How does the transport from the source(s) to the solar system affect their properties? \\item What is the origin of structures observed in their energy distributions? \\item Is there an upper end to the energy spectrum?  \\end{enumerate} The enormous interest particularly in the highest energy {\\sc cr}s originates also from the fact that they provide a direct link to the problems of understanding the evolution of the early universe and to physics beyond the standard model of particle physics in an energy range not accessible to man-made accelerators.  The chemical composition of \\crays measured as a function of their energy provides a key to the answers of above questions and large efforts have been undertaken to provide reliable experimental data for comparison with models of \\cray origin.  In fact, direct \\cray measurements performed at energies up to energies of several GeV have provided a great deal of information about the source composition, propagation effects, the `age' of {\\sc cr}s, etc.  However, for \\eas experiments the situation has been proven much harder than expected.  This is mostly because of the large fluctuations of \\eas parameters and because of uncertainties in corresponding \\eas simulations.  Nevertheless, significant progress has been made in recent years and we have reached a point where high quality \\eas data {\\em and} models of \\cray origin and acceleration can be confronted to each other. The information extracted from such measurements may be complemented by high energy (TeV) gamma and neutrino radiation. The latter messengers point directly to the source but are only secondary probes since they originate from interactions of \\crays with the local environment at the source.  Understanding the common picture of photons, neutrinos and high quality \\cray energy and composition data may finally yield a conclusive picture to the open {\\sc cr}-questions of the century. ",
        "conclusions": ""
    },
    "0212/astro-ph0212512_arXiv.txt": {
        "abstract": "{We present the first high angular resolution 1.4\\,mm and 2.7\\,mm continuum maps of the T\\,Tauri binary system HK\\,Tau obtained with the Plateau de Bure Interferometer. The contributions of both components are well disentangled at 1.4\\,mm and the star previously known to host an edge-on circumstellar disk, HK\\,Tau\\,B, is elongated along the disk's major axis. The optically bright primary dominates the thermal emission from the system at both wavelengths, confirming that it also has its own circumstellar disk. Its non-detection in scattered light images indicates that the two disks in this binary system are not parallel. Our data further indicate that the circumprimary disk is probably significantly smaller than the circumsecondary disk. \\\\ We model the millimeter thermal emission from the circumstellar disk surrounding HK\\,Tau\\,B. We show that the disk mass derived from scattered light images cannot reproduce the 1.4\\,mm emission using opacities of the same population of submicron dust grains. However, grain growth alone cannot match all the observed properties of this disk. We propose that this disk contains three separate layers: two thin outer surfaces which contain dust grains that are very similar to those of the ISM, and a disk interior which is relatively massive and/or has experienced limited grain growth with the largest grains significantly smaller than 1\\,mm. Such a structure could naturally result from dust settling in a protoplanetary disk. ",
        "introduction": "Circumstellar dusty disks are primarily studied to constrain the relevant timescales and various phases of the planetary formation process, from submicron interstellar medium (ISM) dust material to km-sized bodies. The dramatic differences between disks surrounding 10--20\\,Myr-old zero-age main sequence stars such as $\\beta$\\,Pic (Lagrange et al. 2000 and references therein) and those around the much younger T\\,Tauri stars strongly points toward an evolutionary process that occurs within $\\lesssim10^7$\\,yrs. T\\,Tauri stars are thus obvious targets for studying the very first stages of disk evolution and, more specifically, the growth of dust grains in protoplanetary disks from ISM-sized particules to significantly larger ones.  The dust component of circumstellar disks can be traced through its thermal continuum emission. This is especially straightforward at millimeter wavelengths, at which the disk is almost entirely optically thin. Such studies have been performed for over a decade and first hinted at the presence of ``large'' dust grains, as evidenced by ``abnormal'' dust emissivity properties (Beckwith et al. 1990, hereafter B90). However, disk masses and maximum grain size are currently poorly constrained due to the uncertain grain shape and composition (Pollack et al. 1994; Henning \\& Stognienko 1996). Furthermore, the low spatial resolution of most radio observations only rarely allows direct studies of the structure of the disks nor do they resolve the tight visual binary systems commonly found among pre-main sequence stars (Duch\\^ene 1999 and references therein). It is also possible to study the dust grain population of circumstellar disks through high contrast, high-angular resolution imaging of scattered light in the environment of pre-main sequence stars. Such observations, however, are only sensitive to those grains whose sizes are within a factor of a few of the wavelength. The presence (or absence) of much smaller/bigger grains is not constrained by scattered light images. Only the combination of the thermal and scattering approaches can fully determine the properties of dust in circumstellar disks.  {\\hk} is a 2\\farcs4 pre-main sequence binary system located in the nearby Taurus star-forming region ($D\\approx140$\\,pc, Bertout et al. 1999). The primary star presents the usual caracteristics of a classical T Tauri star surrounded by a circumstellar accretion disk (Cohen \\& Kuhi 1979, Kenyon \\& Hartmann 1995). However, the secondary component, which was discovered by Cohen \\& Kuhi (1979), is remarkably faint at visible/near-infrared wavelengths. B90 and Beckwith et al. (1991, hereafter BS91) first detected continuum millimeter and submillimeter emission from the system and estimated a first disk mass, $M_d \\sim 10^{-2}\\,M_\\odot$, although their large beam sizes did not resolve the binary.  The {\\hk} system was later observed at high angular resolution in the near-infrared by Koresko (1998, hereafter K98) and by Stapelfeldt et al. (1998, hereafter S98) through deep HST/WFPC2 visible imaging. {\\hka} appeared as a bright point source while {\\hkb}\\footnote{In some early studies this component was named {\\hk}/c. We adopt here the convention proposed by the IAU and name the secondary component {\\hkb}.} offered a dramatically different shape, that of an extended nebulosity, about 210\\,AU across, with a central dark lane. Both interpreted these images as the signature of an optically thick circumstellar disk seen almost edge-on that completely obscures the central star and only lets scattered photons reach the observer. They further estimated that the total disk mass needed to reproduce these images assuming ISM-like dust grains is 1--2 orders of magnitude smaller than that derived by B90 and BS91. More recently, D'Alessio et al. (2001, hereafter DACH01) modeled the disk of {\\hkb}, trying to match simultaneously the millimeter continuum flux of B90 and the morphology of the disk as seen in optical scattered light. They assumed a disk mass several times larger than that estimated by B90 and BS96 and, allowing the grain size distribution to vary, found a reasonable match to the data for a maximum grain size between 1\\,mm and 1\\,m. The large discrepancies with the conclusions of S98 and K98 raised doubts about the actual disk's structure and dust content.  However, to properly estimate the disk mass or to conclude that dust grains have grown from ISM-like particules, it is necessary to remove the ambiguity induced by the {\\it unresolved} millimeter measurement of {\\hk} from B90. In this paper, we present new 1.4\\,mm and 2.7\\,mm images obtained with the IRAM Plateau de Bure Interferometer with an angular resolution that allows us to resolve the binary system at the shortest wavelength. The observations and results are described in \\S\\,\\ref{sec:obs} and \\S\\,\\ref{sec:res} respectively, while a model of the 1.4\\,mm emission is presented in \\S\\,\\ref{sec:model}. In \\S\\,\\ref{sec:discus}, we evaluate the implications of our observations for the dust content of the disk surrounding {\\hkb} as well as for the geometrical properties of the whole system. Finally, we summarize our findings in \\S\\,\\ref{sec:concl}. ",
        "conclusions": "\\label{sec:concl}  We have obtained new high angular resolution continuum maps at 1.4\\,mm and 2.7\\,mm of the pre-main sequence binary system {\\hk}. For the first time, the binary is clearly resolved at 1.4\\,mm and we have set an upper limit on the flux density of the faintest component at 2.7\\,mm. The emission from both objects appears to be thermal and, while the 1.4\\,mm continuum emission from {\\hkb} is resolved along the major axis of its previsouly known edge-on circumstellar disk, {\\hka} is found to be unresolved at both wavelengths. These observations confirm that both stars are surrounded by their own circumstellar disk. We further find that the disk surrounding {\\hka} must be observed at a less favorable inclination ($i\\lesssim70\\degr$) than that of its companion, which could be the result of turbulence-induced fragmentation, while its radius ($\\lesssim60$\\,AU) is unexpectedly smaller than that of the circumsecondary disk.  Considering the various observational constraints on the circumstellar disk of {\\hkb}, we suggest that it must have a layered structure, such as investigated by Chiang et al. (2001). The upper/lower layers, which are displaced vertically from the disk midplane, do not contain a significant population of grains larger than 1\\,$\\mu$m in radius and thus have dust properties very similar to those of the ISM. On the other hand, the interior of the disk is either massive and constituted of ISM-like grains or has a population of grains that extends significantly beyond $\\mu$m-sized particules. However, the grain size distribution cannot extend to 1\\,mm or beyond. Such a disk structure, although suggested for another system, is not required for most disks around T\\,Tauri stars but could naturaly result from dust settling on very short timescales.  Only future high-angular observations of the system covering the mid- and far-infrared ranges will allow a complete study of the disk structure and dust content of {\\hkb}, via the analysis of its complete SED. For instance, measurements of the 10--20\\,$\\mu$m flux from {\\hkb}, where the SED is close to its minimum between the scattering and thermal regimes, would be extremely valuable. NGST and ALMA are among the scheduled new facilities that will allow such observations. A more detailed study of the dust population at the surface of the disk will also be possible through the analysis of visible/near-infrared polarization maps of the disk."
    },
    "0212/astro-ph0212038_arXiv.txt": {
        "abstract": "We report the discovery of H$_2$O maser emission at 1.35\\,cm wavelength in seven  active galactic nuclei (at distances up to $<80$ Mpc) during a survey conducted at the 70-m diameter antenna of the NASA Deep Space Network near Canberra, Australia.  The detection rate was $\\sim 4\\%$.  Two of the maser sources are particularly interesting because they display satellite high-velocity emission lines, which are a signature of  emission from the accretion disks of supermassive black holes when seen edge on.  Three of the masers are coincident, to within uncertainties of $0\\rlap{.}''2$, with continuum emission sources we observed at about $\\lambda 1.3$~cm. We also report the discovery of new spectral features in the Circinus galaxy H$_2$O maser that broaden the known velocity range of emission therein by a factor of $\\sim 1.7$. If the new spectral features originate in the Circinus accretion disk, then molecular material must survive at radii $\\sim 3$ times smaller than had been believed previously ($\\sim 0.03$~pc or $\\sim 2\\times10^5$ Schwarzschild radii). ",
        "introduction": "Water maser emission ($\\lambda 1.35$~cm in the rest frame) is known to trace warm, dense gas at radii of 0.1 to 1 pc in the accretion disks surrounding supermassive black holes in galactic nuclei \\citep[e.g.,][]{mgh99}.  It can also trace material heated by jet activity \\citep[e.g.,][]{claussen98} and wide-angle nuclear winds \\citep{greenhill.circinus}. Emission from disks is visible when they are viewed close to edge-on and amplification paths are longest.  Several have been mapped with Very Long Baseline Interferometry (VLBI): NGC\\,4258 \\citep{miyoshi95}, NGC\\,1068 \\citep{gg97}, and the Circinus Galaxy \\citep{greenhill.circinus}. ``Maser disks'' may also exist in IC\\,2560 \\citep{ishihara01}, NGC\\,5793 \\citep{hagiwara01}, and  Mrk\\,1419 \\citep{henkel02}, though confirmation awaits further study.  Triply peaked spectra characterize emission from accretion disks that are well populated by masers.  Emission close to the systemic velocity of the host galaxies (i.e., low-velocity emission) occurs where orbital motion is transverse to the line of sight. High-velocity emission is symmetrically offset by the disk orbital velocities and arises in regions where the disk motion is parallel to the line of sight.   Velocities as high as $\\sim 1100$\\kms~have been observed (i.e., in NGC\\,4258).  Water maser sources in active galactic nuclei (AGN) are important astrophysical tracers in part because VLBI can provide maps of resolved disk structure (e.g., warping) and dynamics (e.g., rotation curves and proper motions), as in \\citet{herrnstein99}. Unfortunately, only $\\sim 30$ H$_2$O masers are known in AGN, and few of these exhibit triply peaked spectra.  The discovery of new masers is a priority and a challenge.  First, the emission is typically weak, and surveys must invest substantial time observing each target with the largest available apertures (e.g., 1 hour with a 100 m diameter antenna). Second, mean detection rates in surveys are typically $\\ll 10$\\% for Seyfert~II galaxies closer than $cz\\sim 7000 $\\kms~\\citep{braatz97}.  Third, because the orbital speeds of disks (and concomitant velocity range of maser emission) cannot be known in advance, surveys must have instantaneous bandwidths of thousands of \\kms. Sufficiently broadband observing systems have become available to the general community only recently.  We report the detection of seven new masers obtained in a high sensitivity, broad bandwidth survey of 160 nearby AGN ($cz<8100$\\kms).  The survey was unusual for two reasons.  First, we used a 70-m antenna of the NASA Deep Space Network (DSN) to achieve high  sensitivity \\citep[see also][]{greenhill3735}. Second, we used a custom built, portable, 5350\\kms-wide spectrometer.  The survey, which is ongoing at DSN facilities in the Northern and Southern Hemispheres, and the hardware will be discussed in detail elsewhere. ",
        "conclusions": "We have detected new H$_2$O maser sources in seven AGN, as well as new high-velocity emission components in the Circinus galaxy.  Two of the new masers exhibit high-velocity emission, indicative of emission from accretion disks.  It should be possible to map these two disks with VLBI, trace their rotation curves, and weigh the central engines that bind them.    Ultimately, it may be possible to measure geometric distances to the two host galaxies, which would be significant because both lie in the Hubble Flow.  One of the new masers lies in an early-type galaxy, and we speculate that it might therefore be associated with an as yet undetected jet.  We have also discovered new high-velocity water maser emission in the Circinus galaxy, VLBI observation of which could greatly improve our understanding of its accretion disk."
    },
    "0212/astro-ph0212042_arXiv.txt": {
        "abstract": "Stellar evolution calculations predict the flux-weighted gravity $g$/\\teffq\\/ and absolute bolometric magnitude of blue supergiants to be strongly correlated. We use medium resolution multi-object spectroscopy of late B and early A supergiants in two spiral galaxies, NGC 300 and NGC 3621, to demonstrate the existence of such a relationship, which proves to be surprisingly tight. An analysis of high resolution spectra of blue supergiants in Local Group galaxies confirms this detection. We discuss the application of the relationship for extragalactic distance determinations and conservatively conclude that once properly calibrated it has the potential to allow for measurements of distance moduli out to 30.5 mag with an accuracy of 0.1 mag or better. ",
        "introduction": "Bright supergiants in external galaxies have been recognized as valuable extragalactic distance indicators for a long time, since the pioneering work of \\citet{hubble36}. A number of works have attempted to calibrate photometric and/or spectroscopic signatures of blue supergiants (those of spectral type OBA) for this purpose, but the uncertainty in the derived distances (typically 0.4 mag or larger in distance modulus) have always been a major drawback for these techniques (\\citealt{humphreys88}, \\citealt{tully84}).  The discovery of a wind momentum-luminosity relationship (WLR) for blue supergiants (\\citealt{kud00}) exploits the dependency of the strength of the radiation driven winds of massive stars on stellar luminosity, offering a potentially more accurate distance indicator.  While hot O stars provide so far the best calibration of the WLR (\\citealt{puls96}, \\citealt{puls02}), it is the visually brightest late B and early A supergiants ($M_V\\simeq-9$) which offer the largest potential as extragalactic standard candles (\\citealt{kud98}, \\citealt{kud99}). This can now be investigated for the nearby galaxies ($D<10$ Mpc) with multiobject spectroscopy at 8-m class telescopes, as the exploratory work of \\citet{bresolin01, bresolin02} has shown. Quantitative spectroscopy of individual BA supergiants leads to the determination of gravities, temperatures, metallicities and stellar wind parameters (based on the wind emission in H$\\alpha$), which are then combined to provide distances.  In this Letter we suggest a novel method, based on the absorption strengths of the higher Balmer lines formed in the photosphere. The concept is not entirely new, since a relationship between the equivalent width of H$\\gamma$ and $M_V$ for Galactic and Magellanic Clouds BA supergiants, together with a temperature dependence, has already been discussed by several authors (\\citealt{petrie65}, \\citealt{crampton79}, \\citealt{tully84}, \\citealt{hill86}).  The recent improvements in the modelling of A supergiant atmospheres in NLTE and the development of new diagnostic tools (\\citealt{venn95a,venn95b,venn99,venn01}, \\citealt{przybilla01a,przybilla01b,przybilla02}) allows us to determine stellar parameters with unprecedented accuracy and reliability.  Based on these achievements, we discuss here a promising application on a set of high-to-intermediate resolution spectra obtained for a sample of A supergiants in the Milky Way, the Magellanic Clouds, NGC~300 and NGC~3621, and few additional objects in a handful of nearby galaxies.  The theoretical concept is explained in \\S 2, and we present the observational tests in \\S 3 and \\S 4. Future work on supergiants in nearby galaxies is briefly summarized in \\S 5. ",
        "conclusions": "The results presented in the previous sections are very encouraging. The flux-weighted gravities of late B and early A supergiants are obviously very tightly correlated with absolute bolometric magnitude. The application of this relationship, once properly calibrated, for extragalactic distance determinations is straightforward. It requires multi-colour photometry of galaxies containing a young stellar population to identify possible blue supergiants and subsequent medium resolution ($\\sim 5\\,$\\AA) multi-object spectroscopy (see \\citealt{bresolin01,bresolin02}) to determine effective temperature and gravity directly from the spectra. The spectral analysis will also yield bolometric correction (which is small for these spectral types) and intrinsic colour so that an accurate correction for reddening and extinction is possible. Application of the FGLR will then provide the absolute bolometric magnitude, which by comparison with the de-reddened visual magnitude will give the distance modulus. Assuming a residual scatter of $\\sigma$ = 0.3 mag for the FGLR (see Fig.~2 and 3) we estimate that with 10 supergiant stars per galaxy we can achieve an accuracy of $0.1$ mag in distance modulus. We are confident that in one night of observing time we can reach down to $V = 22.5$ with the existing very efficient medium resolution multi-object spectrographs attached to 8m-class telescopes. With objects in an absolute magnitude range between $-8$ and $-10$ mag the FGLR method appears to be applicable out to distance moduli of $m-M = 30.5$ or even beyond.  The restriction to medium resolution spectroscopy in the blue spectral range provides significant advantages. Most importantly, contamination from sky and \\hii~region emission is by far less critical than in the red, which is needed as an additional spectral range, for instance, for the WLR method, which requires the measurements of H$\\alpha$ profiles with at least $\\sim2\\,$\\AA\\/ resolution.  Moreover, the amount of observing time for an accurate distance determination is significantly reduced, if only spectra in the blue are required.  The accurate calibration of the FGLR (and the WLR) will become the crucial element of future work, before the method can be applied seriously for extragalactic distance determinations. Local Group galaxies with well determined distances provide the ideal laboratory for this purpose. Multi-object spectrographs with rather high spectral resolution attached to 8 to 10m-class telescopes such as FLAMES (VLT) or DEIMOS (Keck 2) will allow high quality spectra of large candidate samples in each galaxy with a rather modest amount of observing time.  Such a systematic study of hundreds of blue supergiants in Local Group galaxies will not only provide an accurate calibration of the FGLR. It will also enable us to investigate important aspects of stellar evolution, which are related. The most crucial ones concern the role of metallicity and stellar rotation. Investigating stellar evolutionary tracks at different metallicity and with different initial rotation at the main sequence (\\citealt{meynet00,meynet94}) we find small, but noticeable effects on the theoretical FGLR. Mass-loss in evolutionary stages prior to the blue supergiant phase depends on metallicity (\\citealt{kud00}) and has a small influence on the FGLR. Rotation affects the strength of mass-loss and the internal mixing processes and introduces a modification of the mass-luminosity relationship in the blue supergiant stage. The effect of rotation becomes larger at higher luminosities. This might be the reason for the increase of the residual scatter at luminosities above $M_{bol}=-8$ mentioned above.  Another very important issue is the fraction among the sample of observed blue supergiants evolving backwards to the blue after a previous phase as red supergiants. Those objects are expected to have lost a significant fraction of their mass as red supergiants and might form an additional sequence below the observed relationship. Evolutionary calculations indicate that the relative number of those objects might depend crucially on metallicity and rotation. Systematic studies of Local Group galaxies at different metallicity, as proposed above, will allow to investigate this problem.  In general, the observational detection of the tight relationship between flux-weighted gravity and absolute bolometric luminosity is a triumph of two classical areas of astrophysics, stellar evolution and stellar atmospheres. It confirms the general scenario of stellar evolution with mass-loss and rotation away from the main sequence and the predicted mass-luminosity relation. It also confirms the power and accuracy of present day spectroscopic stellar diagnostics. It is very satisfying to see the potential of these disciplines for a significant contribution to quantitative extragalactic studies.   \\clearpage"
    },
    "0212/gr-qc0212005_arXiv.txt": {
        "abstract": " ",
        "introduction": "One of the most intriguing  prediction of the General Theory of Relativity, in its linearized weak--field and slow--motion approximation, is the so called frame--dragging or Lense--Thirring effect (Lense and Thirring 1918). It can be thought of as a consequence of a gravitational coupling between the proper angular momentum \\bm J of a central body of mass $M$, which acts as source of the gravitational field, and the angular momentum \\bm s of a particle freely orbiting it. It turns out that the spin \\bm s undergoes a tiny precessional motion (Schiff 1960). The most famous experiment devoted to the measurement of such a gravitomagnetic effect is the Stanford University GP--B mission (Everitt $et\\ al $ 2001) which is scheduled to fly in April 2003.  If we consider the whole orbit of a test particle in its geodesic motion around $M$ as a sort of giant gyroscope, its orbital angular momentum \\bm l undergoes the Lense--Thirring precession, so that the longitude of the ascending node $\\Omega$ and the argument of pericenter $\\omega$ of the orbit of the test particle (Sterne 1960) are affected by small secular precessions\\footnote{In the original paper by Lense and Thirring the longitude of the pericenter $\\varpi=\\Omega+\\omega$ is used instead of $\\omega$. } $\\dot\\Omega_{\\rm LT}$, $\\dot\\omega_{\\rm LT}$ (Lense and Thirring 1918; Ciufolini and Wheeler 1995; Iorio 2001). \\subsection{The LAGEOS-LAGEOS II Lense-Thirring experiment} Up to now, the only attempts to detect them in the gravitational field of the Earth are due to Ciufolini and coworkers (Ciufolini $et\\ al$ 1998; Ciufolini 2000; 2002) who analysed the laser data of the existing geodetic SLR (Satellite Laser Ranging) satellites LAGEOS and LAGEOS II over time spans of some years. The observable is a suitable combination of the orbital residuals of the nodes of LAGEOS and LAGEOS II and the perigee of LAGEOS II according to an idea exposed in (Ciufolini 1996). The relativistic signal is a linear trend with a slope of almost 60.2 milliarcseconds per year (\\msy\\ in the following). The claimed total accuracy is of the order of $20\\%-30\\%$. The main sources of systematical errors in such kind of measurements are the unavoidable aliasing effect due to the mismodelling in the classical secular precessions induced by the even zonal coefficients of the multipolar expansion of the Earth's gravitational field (Kaula 1966) and the non--gravitational perturbations affecting especially the perigee of LAGEOS II. It turns out that the mismodelled classical precessions due to the first two even zonal harmonics of the geopotential $J_2$ and $J_4$ are the most insidious source of error for the Lense--Thirring measurement with LAGEOS and LAGEOS II. The combination of (Ciufolini 1996) is insensitive just to $J_2$ and $J_4$. According to the full covariance matrix of the EGM96 gravity model (Lemoine $et\\ al$ 1998), the error due to the remaining uncancelled even zonal harmonics amounts to almost 13$\\%$. A reliable evaluation of the impact of the geopotential is a particularly subtle and important topic. Indeed, it is based on the use of the covariance matrix of the even zonal harmonics of the geopotential of some terrestrial gravity models like EGM96. As pointed out in (Ries $et\\ al$ 1998), in obtaining the solution of EGM96 and of other previous gravity models a multidecadal observational time span has been used and many seasonal, stochastic and secular variations, which is known that they affect the geopotential, have not been accounted for. Then, according to the remarks of (Ries $et\\ al$ 1998), nothing would assure that during any particular relatively short time span as that used in the LAGEOS--LAGEOS II Lense--Thirring experiment the correlation among the geopotential coefficients would be just that of the EGM96 covariance matrix. For example, it seems that the $13\\%$ favorable estimate is based on a particular correlation between $J_6$ and $J_8$ (Ries $et \\ al$ 1998) which do affect the observable by Ciufolini. If, with a more conservative, although maybe pessimistic, approach, we use the diagonal part only of the covariance matrix of EGM96 up to degree $l=20$ \\zone for the LAGEOS--LAGEOS II Lense--Thirring experiment amounts to  $46.6\\%$ (Iorio 2002a). Moreover, it turns out to be insensitive to the even zonal harmonics of degree higher than $l=20$ due to the high altitude of the LAGEOS satellites. Indeed, the classical precessions of the node and the perigee depend on $\\left(\\frac{R}{a}\\right)^l a^{-\\rp{3}{2}}$, where $R$ is the Earth's radius and $a$ is the satellite's semimajor axis. \\subsection{The LARES project} The originally proposed LARES mission (Ciufolini 1986; 1998) consists of the launch of a LAGEOS--type satellite--the LARES--with the same orbit of LAGEOS except for the inclination $i$ of its orbit, which should be supplementary to that of LAGEOS, and the eccentricity $e$, which should be one order of magnitude larger. In Table 1 the orbital parameters of the existing and proposed LAGEOS--type satellites are quoted. \\begin{table}[ht!] \\caption{Orbital parameters of \\lg, \\lgg, LARES and POLARES and their Lense--Thirring precessions.} \\label{tavola1} \\begin{center} \\begin{tabular}{llllll} \\noalign{\\hrule height 1.5pt} Orbital parameter & \\lg & \\lgg & LARES & POLARES\\\\ \\hline $a$ semi major axis (km) & 12,270 & 12,163 & 12,270 & 8,378\\\\ $e$ eccentricity & 0.0045 & 0.014 & 0.04 & 0.04\\\\ $i$ inclination (deg) & 110 & 52.65 & 70 & 90\\\\ $\\dot\\Omega_{\\rm LT}$ (\\msy) & 31 & 31.5 & 31 & 96.9\\\\ $\\dot\\omega_{\\rm LT}$ (\\msy)& 31.6 & -57 & -31.6 & 0\\\\ \\noalign{\\hrule height 1.5pt} \\end{tabular} \\end{center} \\end{table} The choice of the particular value of the inclination for the LARES is motivated by the fact that in this way, by using as observable the sum of the nodes of the LAGEOS and the LARES, it should be possible to cancel out exactly all the contributions of the even zonal harmonics of the geopotential, which depends on $\\cos i$, and add up the Lense--Thirring precessions which, instead, are independent of $i$.  Of course, it would not be possible to obtain practically two orbital planes exactly 180 deg apart due to the unavoidable orbital injection errors. It turns out that all depends on the last stadium of the rocket used. According to conservative estimates, if a solid propellant is used for it an error in inclination of the order of 1 deg is to be expected, while if a liquid propellant, which is more expensive, is used  the error should amount to 0.5-0.6 deg (Anselmo, private communication 2002). In Figure 1, page 4314 of (Iorio $et \\ al$ 2002) the impact of such source of error on the originally proposed LAGEOS--LARES mission has been shown. It should be noted that the simple sum of the nodes of the LAGEOS and the LARES is affected by all the even zonal harmonics of the geopotential. Then, when the impact of the departures of the inclination of LARES from its nominal values on \\zone has to be calculated, the role of all the correlations among the even zonal harmonics should be important. If we decide to take into account the remarks of (Ries $et\\ al$ 1998), the results obtained in (Iorio $et\\ al$ 2002) might be considered optimistic in the sense that they are based on an extrapolation of the validity of the full covariance matrix of EGM96 up to degree $l=20$ to arbitrary future time spans.  In (Iorio $et\\ al$ 2002) an alternative observable based on the combination of the residuals of the nodes of LAGEOS, LAGEOS II and LARES and the perigee of LAGEOS II and LARES has been proposed. It would allow to cancel out the first four even zonal harmonics $J_2,\\ J_4,\\ J_6,\\ J_8$ so that \\zone would be rather insensitive to the orbital injection errors in the LARES inclination and would amount to $0.02\\%$ only.  In (Iorio 2002b) the recent proposal of inserting the LARES in a low--altitude polar orbit (Lucchesi and Paolozzi 2001), so to obtain the so called POLARES, and to analyze only its node has been critically analyzed from the point of view of the impact of the orbital injection errors in the POLARES inclination. \\subsection{Motivation of the present work} The conclusions obtained in (Iorio $et\\ al$ 2002; Iorio 2002b) are based on the assumption of the validity of the EGM96 full covariance matrix in arbitrary future time spans during which, instead, it might happen that the correlations between the even zonal harmonics will be different. This problem is particularly relevant for those observables which are sensitive to the full range of the even zonal harmonics of the geopotential like the originally proposed node--only LAGEOS--LARES combination and the node--only POLARES observable.  Consequently, we wish to reanalyze such issues in a more conservative, although pessimistic, approach by using the diagonal part only of the covariance matrix of EGM96. We expect that \\zone of the new proposed observable based on the use of the orbital elements of LAGEOS, LAGEOS II and LARES, which is $J_2-J_8$-free, should be relatively insensitive to the correlation among the remaining even zonal harmonics, contrary to the sum of the nodes of the LAGEOS and the LARES and the node of the POLARES. If it will be so, the reliability of the modified version of the LARES project will be enforced and posed on a more firm basis. ",
        "conclusions": "In this letter we have followed a more conservative and realistic approach in evaluating the impact of the orbital injection errors of LARES and POLARES on \\zone of some proposed observables aimed to the measurement of the Lense--Thirring effect with LAGEOS--like SLR satellites. In particular, we have used, in a root--sum--square fashion, the diagonal part only of the covariance matrix of the even zonal harmonic coefficients of the EGM96 Earth's gravity model up to degree $l=20$.  With regard to the LARES project, it turns out that the recently proposed residuals combination involving the LAGEOS, LAGEOS II and LARES satellites, which cancels out the first four even zonal harmonics, should yield not only a more accurate measurement than the simple sum of the nodes of LAGEOS and LARES, but also more reliable because it would be less dependent on the correlation between the remaining even zonal harmonics of higher degree.  The use of POLARES, according to the present level of knowledge of the terrestrial gravitational field and to the results presented here, should be considered unpracticable. It is probable that, even with the new gravity models from CHAMP and GRACE, such a proposed configuration would not be competitive with the other multisatellite observables which, in turn, would benefit of the new data of the gravitational field.  The new gravity models from the CHAMP and GRACE missions should yield great benefits for a more confident and reliable assessment of \\zone in all the examined missions. However, also in this case the originally proposed LARES orbital configuration, in conjuction with the data from LAGEOS and LAGEOS II, should be the best choice."
    },
    "0212/astro-ph0212274_arXiv.txt": {
        "abstract": "We use a simple statistical test to show that the anomalous flux ratios observed in gravitational lenses are created by gravitational perturbations from substructure rather than propagation effects in the interstellar medium or incomplete models for the gravitational potential of the lens galaxy. We review current estimates that the substructure represents $0.006 < f_{sat} < 0.07$ (90\\% confidence) of the lens galaxy mass, and outline future observational programs which can improve the results. ",
        "introduction": "It is a generic feature of CDM (cold dark matter) halo simulations that a significant fraction of the halo mass remains in the form of satellites (e.g. Kauffmann et al.~1993, Moore et al.~1999, Klypin et al.~1999).  The exact mass fraction remains somewhat unclear, but the global mass fraction is of order 5--10\\%, and the projected fraction inside cylinders of radius $R\\sim 5h^{-1}$kpc is of order 1\\%.  These mass fractions are significantly higher than are observed in satellites of the Galaxy, suggesting a conflict between CDM models and observations.  Three general classes of solutions have been proposed.  First, the satellites can be made invisible by suppressing star formation (Kauffmann et al.~1993; Bullock et al.~2000). Second, they can be destroyed by normal dynamical processes or abnormal ones such as self-interacting dark matter (Spergel \\& Steinhardt~2000; Yoshida et al.~2000; Col{\\'\\i}n et al.~2002; D'Onghia \\& Burkert~2002).  Third, their formation might be avoided by significantly reducing the amplitude of the power spectrum on the relevant scales (Kamionkowski \\& Liddle~2000; Col{\\'{\\i}}n et al.~2000; Bode et al.~2001; Avila-Reese et al.~2001). The problem, of course, is that it is difficult to distinguish between undetectable and absent satellites.  It was realized early in the debate (see Moore et al.~1999) that gravitational lensing provided a means of resolving the issue because it could detect the gravitational perturbations created by substructure. It was already known that satellites provided a means of solving the ``anomalous flux ratio'' problem seen in some gravitational lenses (Mao \\& Schneider~1998).  An example of such a problem is shown in Fig.~\\ref{fig:ism}, where the close image pair would be expected to have very similar fluxes for any lens model where the gravitational potential can be well-represented by a low-order Taylor series expansion near the images. Metcalf \\& Madau~(2001) and Chiba~(2002) pointed out that such anomalies should be common given the predicted CDM substructure fractions. Mao \\& Schneider~(1998), Keeton~(2001), Bradac et al.~(2002), and Chiba~(2002) explored how substructure could explain the anomalous flux ratios in several lens systems. Metcalf \\& Zhao~(2001), Keeton, Gaudi \\& Petters~(2002) and Evans \\& Witt~(2002) explored whether the model for the primary lens galaxy could be modified to explain the anomalous flux ratios, with mixed results which we will discuss in detail below.  Our contribution in Dalal \\& Kochanek (2002, DK02 hereafter) was to analyze the data to make an experimental determination of the substructure fraction, finding it to be in the range $0.006 < f_{sat} < 0.07$ (90\\% confidence) based on a sample of radio lenses.  This is in good agreement with the expectations for CDM, and well above standard estimates for the mass fractions in normal satellites.  \\begin{figure}  \\includegraphics[height=.4\\textheight]{fig1.ps} \\includegraphics[height=.4\\textheight]{fig2.ps}  \\caption{(LEFT) Example of an anomalous flux ratio.  For a smooth potential we would expect the A and B images in B1555+375 to have the same flux (Marlow et al.~1999). \\newline (RIGHT) ISM properties needed to explain anomalous flux ratios.  For an optical depth function $\\tau = \\tau_5 (\\nu/5\\hbox{GHz})^\\alpha$ we show the estimated spectral index $\\alpha$ as a function of the optical depth $\\tau_5$ at 5~GHz for the radio lenses used in DK02 with published flux ratios at both 5~GHz and either 8 or 15~GHz. The points are coded by the image type: minima--squares, saddle points--triangles, brightest--filled, and faintest--open. } \\label{fig:ism} \\end{figure}   Our review of the problem will cover three basic topics.  First, we will discuss the problem of distinguishing CDM substructure from other possible origins for the anomalous flux ratios.  In particular, we introduce a simple statistical test for substructure in the gravity as compared to either propagation effects in the interstellar medium of the lens galaxy or poorly modeled contributions to the smooth gravitational potential of the lens galaxy.  In \\S3 we review our estimate of the substructure mass fraction and its relation to simulations.  Finally, in \\S4 we discuss the future of the method. ",
        "conclusions": ""
    },
    "0212/astro-ph0212104_arXiv.txt": {
        "abstract": "Groups and clusters contain a large fraction of hot gas which emits X-ray radiation. This gas yields information on the dynamical state and on the total mass of these systems. X-ray spectra show that heavy elements are present in the gas. As these metals must have been produced in the cluster/group galaxies and later transported into the gas, the metallicity is a good tracer for the transport processes. Several possible processes, that transport gas from the small potential wells of the galaxies into the clusters and groups, are discussed. ",
        "introduction": "Clusters and groups of galaxies do not only contain galaxies but all the space between the galaxies is filled with hot gas. This gas is also often referred to as intra-cluster medium (ICM) or intra-group medium (IGM). The high temperature of the gas of 1 -10 keV is in good correspondence with the depth of the potential wells of groups and clusters (see Table 1). According to the shallower potential of groups also the temperature is lower in groups ($\\approx$ 1 keV) compared to clusters (3-10 keV). The gas is almost fully ionised and has very low densities of $10^{-2} - 10^{-4}$ cm$^{-3}$ decreasing outwards. It is optically thin. In the spectra of both, groups and clusters, lines of heavy elements have been detected corresponding to metallicities of about 0.2 to 0.4 in solar units. These metals indicate that the gas cannot be of purely primordial origin, but part of it must have been previously the inter-stellar medium (ISM) in galaxies and then it must have been transported from the galaxies into the ICM and IGM, respectively, by certain processes. The processes will be discussed in the next section.  \\begin{table} % \\begin{tabular}{lll} \\hline & Clusters  & Groups\\\\ \\hline Temperature   & 3-10 keV       & $\\approx 1$ keV\\\\ Extent        & few Mpc        & 0.1-1 Mpc\\\\ Metallicity   & 0.2-0.4 solar  & 0.2 solar\\\\ X-ray luminosity & $10^{43-45}$ erg/s & $10^{42-43}$ erg/s\\\\ \\hline \\end{tabular} \\caption[]{Properties of the gas in cluster and groups of galaxies.} \\label{parset} \\end{table}  Gas with such properties emits thermal bremsstrahlung in the X-ray range (see Table 1). Therefore groups and clusters are currently studied extensively by the X-ray observatories XMM and CHANDRA. While in almost all of the clusters X-ray emission has been found, the X-ray emission from groups is somewhat harder to detect due to their lower luminosity. So far in about 50\\% of the nearby groups X-ray emission has been found.  \\begin{figure} % \\psfig{figure=schindler_fig1_gimp.ps,width=8.2cm,clip=} \\caption[]{CHANDRA image of the cluster RBS797 (from Schindler et al. 2001). The cluster looks well virialised with a very regular surface brightness distribution.} \\label{RBS797} \\end{figure} ",
        "conclusions": "So far the picture of the interaction between ICM/IGM and the ISM is not very clear. There are many different processes possible that can transport gas from the galaxies to the ICM/IGM: ram-pressure stripping, galactic winds, galaxy-galaxy interaction, jets from active galaxies and maybe even more. It is not known yet, which of the processes is most efficient, what the time variations are and how much it depends on other parameters like e.g. the mass of the cluster/group. The best property to test these different processes is the metallicity. The distribution and the time evolution of the metals are now finally measurable with high accuracy with the X-ray satellites XMM and CHANDRA, so that we will soon get answers on the importance of the different enrichment processes."
    },
    "0212/astro-ph0212332_arXiv.txt": {
        "abstract": "We present the results of a spatially-resolved spectroscopic analysis of the galaxy cluster Abell 496 with the S3 chip on-board the {\\sl Chandra} satellite. We confirm the presence of a central positive temperature gradient consistent with a cooling flow, but with a minimum gas temperature of $\\sim$0.5--0.9 keV. The cluster also exhibits sharp edges in gas density and temperature which are consistent with ``cold front'' substructures. The iron abundance profile is not radially symmetric relative to the cluster center. Towards the direction of the most prominent (northerly) cold front, the iron abundance is roughly flat, with nearly solar values. In the opposite (southerly) direction from the center, the iron abundance distribution shows an ``off-center'' peak. Various abundance ratios suggest that the heavy elements in the central regions of the cluster are dominated by SN Ia ejecta. However, for radii greater than 100 $h_{50}^{-1}$ kpc, the abundance ratios vary in such a way that different abundance ratios provide very different estimates of the proportion of SN Ia/II ejecta. Nonetheless, observed abundances and abundance ratios are continuous across the cold fronts, which suggests that the cold fronts are not likely to be the result of a subcluster merger. We suggest instead that the cold fronts in A496 are caused by ``sloshing\" of the central cooling flow gas, induced by the motion of the cD about the cluster center. ",
        "introduction": "{\\sl Chandra} satellite observations have revealed a wealth of X-ray structures in the cores of galaxy clusters, including X-ray plumes (e.g. Sanders \\& Fabian 2002), cavities (e.g. McNamara et al.\\ 2000), and ``cold fronts'' (e.g.\\  Markevitch et al.\\ 2000; Vikhlinin et al.\\ 2001; Mazzotta et al.\\ 2001). Cold fronts (hereafter CF) are sharp surface brightness discontinuities (widths are typically smaller than the electron mean free path) characterized by a jump in gas temperature, accompanied by a fall in X-ray surface brightness such that the gas pressure remains continuous across the front. Thus, these structures differ from bow shocks (e.g. Markevitch et al.\\ 2002) and are often attributed to subsonic (transonic) motions of accreted substructures (Markevitch et al.\\ 2000, 2001; Vikhlinin et al.\\ 2001)  The incidence of gaseous substructures at the centers of clusters suggests that cluster cores are very dynamic, even when there are no obvious signs of merging (such as in cooling flow clusters with regular isophotes). In this {\\it Letter} we show that this is the case for A496. This cluster is a typical, bright, nearby (z$\\approx$0.032), apparently well-behaved cooling flow cluster. In a joint {\\sl Ginga} and  {\\sl Einstein} analysis of A496, White et al.\\ (1994) found evidence for a central metal abundance enhancement. This was corroborated by {\\sl ASCA} (e.g. Dupke \\& White 2000a), {\\sl SAX} (Irwin \\& Bregman 2001) and {\\sl XMM} (Tamura et al.\\ 2001) with greater statistical significance. Furthermore, Dupke \\& White (2000a) also discovered radial gradients in various abundance $ratios$, indicating that the gas in the central 2-3$^\\prime$ has a higher proportion of SN Ia ejecta ($\\sim$70\\%) than the outer parts of the cluster.  Spectral analysis of radially averaged annuli with {\\sl XMM} is roughly consistent with {\\sl ASCA} results (Tamura et al.\\ 2001). The spatial resolution achieved by the instruments on-board {\\sl XMM} are very limited in establishing fine correlations between spatial substructures, such as CFs, and spectral parameters such as metal abundances and temperatures. However, this analysis  is important since it provides % clues about the nature of such structures. For example, if CFs are caused by the passage of a high velocity gas clump through the intracluster gas, the sharp surface discontinuity should be accompanied by chemical discontinuities as well, representing the enrichment history of the two different systems. This kind of analysis can be performed best with {\\sl Chandra}, given its high angular resolution. In this {\\it Letter} we describe the discovery of CFs in A496 and search for chemical discontinuities across them. A more detailed analysis of the gas mass distribution will be published elsewhere. All distances shown are calculated assuming a H$_0=50$ km~s$^{-1}$Mpc$^{-1}$ and $\\Omega_0=1$. At the distance of this cluster $1^{\\prime\\prime}\\approx$ 0.95 kpc. ",
        "conclusions": "Two recently proposed hypotheses for generating CFs in clusters involve the infall of galaxy groups into the larger clusters. In some cases the substructures are sub(/trans)sonic remnants of mergers (e.g. Markevitch et al.\\ 2000; Vikhlinin et al.\\ 2001). In other cases, such as A1795 (Markevitch et al.\\ 2001), it is proposed that the infall of a small gravitational substructure disturbs the central gravitational potential of the main cluster, causing low-entropy gas in the center to oscillate around some equilibrium position (the ``sloshing'' hypothesis). In both cases one expects the CF to be a surface discontinuity between a cold core moving through hotter, diffuse intracluster gas, with inefficient mixing between the different gas phases across the surface discontinuity. If this were true for A496 one should expect to see different metal enrichment histories across the CF. We find instead that abundances and abundance ratios are continuous across CFs, a result that is independent of calibration uncertainties in determining the absolute abundances of individual elements.  The data is roughly consistent with a model where the cD is oscillating around the clusters' potential well. This model is similar to that suggested for A1795 (Fabian et al.\\ 2001). In both clusters the CFs are aligned with the most likely projected axis of motion of the cD. This axis is inferred by analysis of the direction of the main extension of optical filaments (see Fabian et al.\\ 1981), X-ray and optical isophotal elongation. The cDs in both systems have similar peculiar velocities after correcting for substructure (Bird 1994).  In this model the cD would be dragging/smearing SN Ia Fe enriched cold gas within the spatial oscillation length ($\\sim$ 80 kpc), which defines the distance to farthest CF observed. The most likely evolutionary time frame for this model would place the cD in A496 moving south beginning to generate a second CF. The motion of the cD may add to other competing core heating mechanisms, which may keep $kT^{\\rm cf}_{\\rm min}$ $\\sim1$ keV. A detailed analysis of this model and other alternates is beyond the scope of this {\\it Letter} and will be published elsewhere. The investigation of the behavior of abundances and their ratios across SB discontinuities in the cores of other clusters will allow us to determine whether such structures are typically induced externally, via subcluster mergers, or internally, perhaps through sloshing associated with cD motion."
    },
    "0212/astro-ph0212369_arXiv.txt": {
        "abstract": "I point out that an effective upper limit of approximately 20 Gyr (for a Hubble constant of 72 km/s/Mpc) or alternatively on the $H_0$-independent quantity $H_0t_0 < 1.47$, exists on the age of the Universe, essentially independent of the unknown equation of state of the dominant dark energy component in the Universe.   Unless astrophysical constraints on the age of the Universe can convincingly reduce the upper limit  to below this value no useful lower limit on the equation of state parameter $w$ for this component can be obtained.  Direct dating by stars does not provide a useful constraint, but model-dependent cosmological limits from supernovae and the CMB observations may.  For a constant value of $w$, a bound $H_0t_0 < 1.1$ gives a limit $ w> -1.5$ ",
        "introduction": "The realization that some unknown form of energy density associated with otherwise empty space appears to dominate the gravitational dynamics of the Universe has changed virtually everything in cosmology.   For example, the future evolution of the Universe becomes largely independent of its geometry (\\cite{kraussturn}), so that a closed universe can expand forever, and an open universe can ultimately collapse.  One of the earliest motivations for assuming the existence a cosmological constant,  the simplest form of dark energy, involved a comparison of the Hubble age of the Universe---determined for an assumed cosmological model on the basis of the observed expansion rate today---with a lower limit on the age of the oldest objects in our galaxy.   In order to resolve the paradox when the latter exceeded the former (\\cite{demarque}), cosmologists were driven to consider the possibility of exotic cosmological models that might allow an older universe for a fixed value of the Hubble constant today.  An accelerating universe allows for this possibility simply because galaxies that are now located at a certain distance from us, and which are moving at some fixed velocity were separating from us at a smaller velocity at earlier times, and thus would have required longer to achieve their present separation than would otherwise be required.   For a flat Universe, in the limit where a cosmological constant dominates the energy density, the Hubble age can approach infinity for any value of the Hubble constant.  We currently have no fundamental understanding of the nature of the dark energy (or perhaps more accurately dark \"pressure\").  It is significant that a lower limit on the age of the Universe determined from globular cluster dating techniques now provides independent evidence for the  existence of dark energy, and puts a limit on the equation of state parameter $wp/\\rho$ ( i.e. w=pressure/energy density) of the dark energy $w <-0.4$ (\\cite{krausschab}).  The question also arises however, if exotic equations of state for dark energy can increase the Hubble age, can one put useful constraints of such equations of state from an {\\it upper} limit on the age of globular clusters, or from direct estimates of the Hubble age itself.  I investigate these questions here. ",
        "conclusions": ""
    },
    "0212/astro-ph0212475_arXiv.txt": {
        "abstract": "{ Recent intensive observations of the BL Lac object OJ 287 raise a lot of questions on the models of binary black holes, processing jets, rotating helical jets and thermal instability of slim accretion disks. After carefully analyzing their radio flux and polarization data, Valtaoja et al. (\\cite{valtaoja00}) propose a new binary model. Based on the black hole mass of $4 \\times 10^8 {\\rm M_\\odot}$ estimated with the tight correlations of the black hole masses and the bulge luminosity or central velocity dispersion of host galaxies, we computed the physical parameters of the new binary scenario. The impact of the secondary on the accretion disk around the primary black hole causes strong shocks propagating inwards and outwards, whose arrival at the jet roots is identified with the rapid increase of optical polarization and the large change of polarization angle at about 0.30 yr after the first main optical flare. An increase of optical polarization, a large rotation of positional angle and a small synchrotron flare at 2007.05 between the optical outbursts at 2006.75 and 2007.89 are expected by the model. With the estimated parameters, we predicated an increase of $\\gamma$-ray flux appearing about 5 days after the first optical/IR peak, which is consistent with the EGRET observations. ",
        "introduction": "OJ 287 (0851+202) at red-shift $z = 0.306$ is one of the most active and best studied BL Lac objects. While it was identified in 1967 (Dickel et al. \\cite{dickel67}), its optical observations were dated back to 1890s, which makes it be one of the best candidates for searching for long-term periodicity. Based on the historical optical light curves, Sillanp\\\"a\\\"a et al. (\\cite{sillanpaa88}) found the large optical outbursts of OJ287 recurrent with a period of 11.65 yr and proposed a binary black hole model, which predicted an outburst in the 1994 fall.  Although the predicted large optical outbursts were detected in 1994 and 1995, it was found that the outbursts are double-peaked and the binary black hole model has to be modified (Sillanp\\\"a\\\"a et al. \\cite{sillanpaa96}). Among the proposed models to explain the periodicity of OJ 287, those of binary black holes (Lehto \\& Valtonen \\cite{lehto96}), procession jet (Katz \\cite{katz97}), rotating helical jets (Villata et al. \\cite{villata98}) and the thermal instability of slim accretion disk (Liu et al. \\cite{liu95}) can explain the periodicity in optical band well. However, radio and polarization observations show that the first of the two peaks is narrow and thermal, and the second one is broad and synchrotronic (Valtaoja et al. \\cite{valtaoja00}; Pursimo et al. \\cite{pursimo00}). Both flares should be synchrotronic and identical in the processing/helical jet models while thermal and similar to each other in the binary black hole model (Valtaoja et al. \\cite{valtaoja00}). Another difficulty with the binary model is that it requires an extremely massive primary of $M= 1.5\\times 10^{10} {\\rm M_\\odot}$ (Pietil\\\"a \\cite{pietila98}), which is at least one order of magnitude larger than the super-massive black hole masses of other BL Lac objects (Barth et al. \\cite{barth02}; Wu et al. \\cite{wu02}; Falomo et al. \\cite{falomo02}).  In the binary black hole model, the outburst thermal optical/IR photons should be up-scattered into $\\gamma$-ray regime by the relativistic electrons in the jets via inverse Compton scattering process. From the parameters given by Pietil\\\"a (\\cite{pietila98}), the estimated time for photons to travel from the impact site to the jets is about 20 days and implies a delay of about three weeks between the optical/IR outbursts and high energy $\\gamma$-ray radiations. Such a large lag is not consistent with the EGRET $\\gamma$-ray observations (Shrader et al. \\cite{shrader96}).  To resolve the difficulties of the models mentioned above, Valtaoja et al. (\\cite{valtaoja00}) suggest a new binary scenario to explain the observations of radio flux and polarization. In their model, the first thermal flare of the outbursts is caused when a secondary black hole hyper-sonically passes through the accretion disk around the primary, and the second synchrotron flare after a viscous time scale $t_{\\rm vis}$ is from the jet as the perigee passage enhances the accretion rate in the accretion disk, leading to increased jet mass flux and formation of strong shocks down the jet. As for such a scenario a constant period agrees with all the sharp outbursts (the first peak), a precession of the binary orbit and an extreme primary black hole are not needed (Valtaoja et al. \\cite{valtaoja00}). However, as the lack of accurate knowledge of the primary black hole mass, Valtaoja et al. (\\cite{valtaoja00}) did not estimate the parameters of the new system. In this letter, we extend Valtaoja et al's model to explain all the observations and try to estimate the physical parameters of the new scenario.  After a brief summary of the observations in Sec.~\\ref{sec:obs}, we describe the model in Sec.~\\ref{sec:model}. In Sec.~\\ref{sec:detail}, we try to estimate the parameters of the model. Our discussions and conclusions are presented in Sec.~\\ref{sec:dis}. $H_0 = 50 \\, {\\rm km\\, s^{-1}\\, Mpc^{-1}}$ and $q_0 = 0.5$ are used throughout the paper. ",
        "conclusions": "\\label{sec:dis}  Basing on the observations of OJ 287 in multi-frequencies and of the host galaxy, we extended and estimated the parameters of the revised binary model initially suggested by Valtaoja et al. (\\cite{valtaoja00}). In the model, the mass of the primary black hole is $4\\times 10^8 {\\rm M_\\odot}$, which is consistent with the upper limit obtained by Valtaoja et al. (\\cite{valtaoja00}) and about 40 times smaller than what the L\\&V's model has given (Pietil\\\"a \\cite{pietila98}). As the semi-major axis is, respectively, about $30 r_{\\rm G}$ in the L\\&V's model and about $4.1 \\times 10^2 r_{\\rm G}$ in the revised model, the two black holes are about ten times closer in the L\\&V's model than in ours (in units of the Schwarzschild radius). With such a less massive primary black hole and larger semi-major axis, the gravitational radiation is much weaker in the new scenario, though we cannot exactly estimate its value as we have only an upper limit of the companion mass. It can be expected that the lifetime of the binary system is much longer in the new model than that in the L\\&V's which is only about $10^4 \\, {\\rm yr}$.  As we took the constant period of $P = 11.886$ yr as in Kidger (\\cite{kidger00}), the predicted next outbursts (Kidger \\cite{kidger00}) are that the first thermal optical flare is at 2006.75 (1 Oct 2006) and the second synchrotron flare is at about 2007.89 (21 Nov 2007). The predictions are consistent within the uncertainties of the period  with those given by Valtaoja et al. (\\cite{valtaoja00}). In addition, an increase of optical polarization and large rotation of PA at 2007.05 (18 Jan, 2007) and brightening in $\\gamma$-ray on 5 Oct, 2006 (at 2006.76) are expected by our model.  We identify the small X-ray variations observed about one week after the first flare with radiations from the hot gas around the secondary black hole. However, it is possible that the X-ray variability is not relevant to the pericenter passage event and cannot be detected during the 2006-2007 outbursts. Another possibility is that the arrival signal of sound wave traveling from the impact site to the jets is the variation in X-ray continuum instead of optical polarization. If so, an extremely high eccentricity $e = 0.98$ is needed. The orbit of the secondary in such a system is completely embedded in the thick accretion disk and no outburst can be observed. Any binary system with such an extremely-high initial eccentricity would very rapidly reduce its eccentricity to $e \\la 0.8$ due to gravitational radiation and the interaction between the secondary and the accretion disk (cf Artymowicz \\cite{artymowicz92}).  If the traveling time scale of shock waves is as short as $t_{\\rm s} /2$, the eccentricity of the binary system is $e \\simeq 0.87$ and most part of the orbit must be embedded in the accretion disk. No sharp optical outbursts can be detected. If the traveling time scale of the shock waves is longer than $t_{\\rm s}$, e.g. $2 t_{\\rm s}$, the orbit becomes very circular with $e = 0.66$ and the two passages are very close to each other and $R_{\\rm min}/R_{\\rm max} \\simeq 0.21$. The outbursts caused by the apogee passage become significant and detectable. In this case, an intersection of two periodic variabilities of $p = 11.886\\, {\\rm yr}$ should have been observed. One possibility is that the orbit is circular of $a \\simeq 640 r_{\\rm G}$ (Note that the observed period $P_{\\rm obs} = P_{\\rm orbit} /2$) and we have actually observed two equal passages with $R_{\\rm min} \\simeq R_{\\rm max} \\simeq a$. In this case, the traveling time scale of shock waves is about $7.3 t_{\\rm s} \\simeq 2.2 \\, {\\rm yr}$, which implies an unreasonably large $\\alpha \\simeq 2.0$. Our conclusion is that the shock traveling time scale of shock waves is between about $t_{\\rm s}/2$ and $2 t_{\\rm s}$. It is reasonable to assume that the rapid increase of the optical polarization and the large rotation of positional angle after about 0.30 yr following the first flare are caused by the shock waves arriving at the jet foot-points. As the shock waves may increase the mass flow into jet and cause shocks down the jets, a small synchrotron outburst corresponding to the change of optical polarization is also expected.  We have assumed that the orbit of the secondary is roughly coplanar with the accretion disk in the calculation of the eccentricity $e$ and $R_{\\rm max}$. If they are not roughly coplanar with each other, the system would have a much higher eccentricity and closer apogee. This problem has already been discussed before.  We used the self-similar solutions to a stationary and vertically averaged one-dimension ADAF. Two-dimensional numerical simulation (Igumenshchev \\& Abramowicz \\cite{igo00}) showed that an ADAF with higher $\\alpha$ ($\\ga 1$) is convectively unstable. For an ADAF with lower viscosity ($\\alpha \\la 0.1$), no bipolar or large scale circular motion is present and $t_{\\rm vis}$ should be passage-independent. For ADAF with a moderate viscosity $\\alpha \\sim 0.3$, while a self-similar solution is adequate, time-dependent bipolar outflows and large scale circulations are important (Igumenshchev \\& Abramowicz \\cite{igo00}). Thus, a passage-dependent time scale $t_{\\rm vis}$ is not inconsistent with our model. However, to fully understand the time dependences of the inner accretion disk, two or even three dimensional numerical simulations are needed. In our model, a time-dependent $t_{\\rm vis}$ gives  a time-dependent viscous parameter $\\alpha$. Other estimated parameters are independent of $t_{\\rm vis}$."
    },
    "0212/astro-ph0212196_arXiv.txt": {
        "abstract": "Far Ultraviolet Spectroscopic Explorer ($FUSE$) spectra of 22 Galactic halo stars are studied to determine the amount of \\ion{O}{6} in the Galactic halo between $\\sim$0.3 and $\\sim$10 kpc from the Galactic mid-plane. Strong \\ion{O}{6} $\\lambda$1031.93 absorption was detected toward 21 stars, and a reliable 3 $\\sigma$ upper limit was obtained toward HD~97991.  The weaker member of the \\ion{O}{6} doublet at 1037.62~\\AA\\ could be studied toward only six stars because of stellar and interstellar blending problems. The measured logarithmic total column densities vary from 13.65 to 14.57 with $<$log$N$$>$= 14.17$\\pm$0.28 (1$\\sigma$). The observed columns are reasonably consistent with a patchy exponential \\ion{O}{6} distribution with a mid-plane density of 1.7$\\times$10$^{-8}$ cm$^{-3}$ and scale height between 2.3 and 4 kpc. We do not see clear signs of strong high-velocity components in \\ion{O}{6} absorption along the Galactic sight lines, which indicates the general absence of high velocity \\ion{O}{6} within 2-5 kpc of the Galactic mid-plane. This result is in marked contrast to the findings of Sembach {\\it et al}. who reported high velocity \\ion{O}{6} absorption toward $\\sim$60\\% of the complete halo sight lines observed by $FUSE$. The line centroid velocities of the \\ion{O}{6} absorption does not reflect Galactic rotation well. The \\ion{O}{6} velocity dispersions range from 33 to 78 \\kms\\ with an average of $<$$b$$>$= 45$\\pm$11 \\kms\\ (1$\\sigma$). These values are much higher than the value of $\\sim$18 \\kms\\ expected from thermal broadening for gas at $T \\sim$ 3$\\times$10$^5$ K, the temperature at which \\ion{O}{6} is expected to reach its peak abundance in collisional ionization equilibrium. Turbulence, inflow, and outflow must have an effect on the shape of the \\ion{O}{6} profiles. Kinematical comparisons of \\ion{O}{6} with \\ion{Ar}{1} reveal that 8 of 21 sight lines are closely aligned in LSR velocity ($|\\Delta V_{LSR}| \\leq$5 \\kms\\ ), while 9 of 21 exhibit significant velocity differences ($|\\Delta V_{LSR}| \\geq$ 15 \\kms\\ ). This dual behavior may indicate the presence of two different types of \\ion{O}{6}-bearing environments toward the Galactic sight lines. The correlation between the \\ion{H}{1} and \\ion{O}{6} intermediate velocity absorption is poor. We could identify the known \\ion{H}{1} intermediate velocity components in the \\ion{Ar}{1} absorption but not in the \\ion{O}{6} absorption in most cases. Comparison of \\ion{O}{6} with other highly-ionized species suggests that the high ions are produced primarily by cooling hot gas in the Galactic fountain flow, and that turbulent mixing also has a significant contribution. The role of turbulent mixing varies from negligible to dominant. It is most important toward sight lines that sample supernova remnants like Loop I and IV. The average N(\\ion{C}{4})/N(\\ion{O}{6}) ratios for the nearby halo (this work) and complete halo (Savage {\\it et al}.) are similar ($\\sim$0.6), but the dispersion is larger in the sample of nearby halo sight lines. We are able to show that the \\ion{O}{6} enhancement toward the Galactic center region that was observed in the $FUSE$ survey of complete halo sight lines (Savage {\\it et al}.) is likely associated with processes occurring near the Galactic center by comparing the observations toward the nearby HD~177566 sight line to those toward extragalactic targets. ",
        "introduction": "\\label{section:intro}  The existence of a hot, low density Galactic halo was theoretically postulated to provide pressure support for cool clouds located at large distances from the Galactic plane (Spitzer 1956). Since that prediction, observational evidence has accumulated in support of this hot phase of the interstellar medium (ISM).  Detections of absorption lines of highly-ionized atoms, such as \\ion{O}{6}, \\ion{N}{5}, \\ion{C}{4}, and \\ion{Si}{4} (York 1974, 1977; Jenkins 1978a; Savage \\& Massa 1987; Sembach \\& Savage 1992; Savage, Sembach, \\& Lu 1997; Savage et al. 2000; Savage {\\it et al}. 2002; Sembach {\\it et al}. 2002; Wakker {\\it et al}. 2002) suggest the widespread, but patchy, distribution of high temperature gas in the Galactic disk and halo. The detections of \\ion{O}{6} are particularly important because the energy needed to ionize \\ion{O}{5} is 113.9~eV, well above the 54.4~eV ionization potential of \\ion{He}{2}. The amount of stellar flux capable of ionizing \\ion{O}{5} in the ISM is limited by strong photospheric \\ion{He}{2} absorption; therefore, it is likely that collisional ionization and not photoionization is the dominant source of \\ion{O}{6}.  The presence of \\ion{O}{6} in the Galactic halo and the detection of soft X-ray emission at high Galactic latitudes (Levine et al. 1976, 1977; Snowden et al. 1998) provide support to the existence of a hot Galactic halo.  Our best access to the interstellar \\ion{O}{6}, the \\ion{O}{6} 1031.93, 1037.62 \\AA\\ resonance doublet, lies in a wavelength range that is only now being explored in detail.   Apart from sporadic and short-lived programs to observe \\ion{O}{6} with the Hopkins Ultraviolet Telescope (Davidsen 1993) and the {\\it Orbiting and Retrievable Far and Extreme Ultraviolet Spectrometers} (Hurwitz \\& Bowyer 1995; Hurwitz et al. 1998; Windmann et al. 1998; Sembach, Savage, \\& Hurwitz 1999), our understanding of the Galactic distribution of this ion was limited until recently to the interstellar \\ion{O}{6} survey by the $Copernicus$ satellite (Jenkins 1978a,b). The situation has changed rapidly with the launch of the $FUSE$ satellite, which was specifically designed for high resolution and high sensitivity studies of the 905 to 1187~\\AA\\ spectral range.  A primary objective of the $FUSE$ program has been to quantify the distribution and properties of hot gas traced by \\ion{O}{6} in the Galaxy and nearby universe. Initial studies of the hot gas in the Magellanic Clouds (Howk {\\it et al}. 2002b; Hoopes {\\it et al}. 2002) and the low-redshift universe (Sembach {\\it et al}. 2001) have already been completed.  A large number of Galactic ($\\sim$200) and extragalactic ($\\sim$100) targets have been observed by the $FUSE$ Science Team for the purpose of mapping the \\ion{O}{6} distribution in the Galactic halo and disk. The halo study is based primarily on a survey of extragalactic objects (Savage {\\it et al}. 2002; Sembach {\\it et al}. 2002), while a sample of early-type stars is being used to study the properties of \\ion{O}{6} in the Galactic disk at distances greater than those accessible with Copernicus (Jenkins 1978a,b). Extragalactic objects -AGNs and QSOs- offer better sky coverage above $|b| \\sim$~20$^{\\circ}$ in the Galactic halo, because only a handful of suitable halo stars are available. Also, observations toward extragalactic objects sample the full extent of halo gas, providing information on the outer regions of the halo where suitable Galactic targets are rare. However, the kinematical structure of the \\ion{O}{6} absorption can be confusing, and it is often difficult to assess where the absorption occurs in the halo. The purpose of including early-type halo stars in the original sample was to provide information about the \\ion{O}{6} distribution at intermediate distances away from the Galactic plane ($|z| \\sim$ 0.5-3~kpc). In the present work, we attempt to provide this information and fill in the distance gap between the disk and halo surveys.  The $FUSE$ sample of Galactic halo stars is also well suited for comparisons of the highly-ionized atoms in regions of the Galaxy where many different physical processes may be occurring (see Shull \\& Slavin 1994). Most of the stars in our sample have been observed previously with the International Ultraviolet Explorer ($IUE$), the Goddard High Resolution Spectrograph (GHRS), or with the Space Telescope Imaging Spectrograph (STIS). Column density measurements for \\ion{Si}{4}, \\ion{C}{4}, and \\ion{N}{5} are readily available in the literature (Sembach \\& Savage 1992, 1994; Savage \\& Sembach 1994; Savage, Meade, \\& Sembach 2001; Savage, Sembach, \\& Howk 2001) for comparison with \\ion{O}{6}. The comparisons of the \\ion{Si}{4}, \\ion{C}{4}, \\ion{N}{5}, and \\ion{O}{6} columns are important because the models of the hot gas in the Galactic halo predict different density ratios for the highly-ionized species. Several physical processes have been proposed to explain the production of these species in or near the interfaces between the hot and warm gas. These are the cooling Galactic fountain models (Shapiro \\& Field 1976; Shapiro \\&  Benjamin 1991), cooling supernova remnant (SNR) models (Slavin \\& Cox 1992, 1993; Shelton 1998), and models that involve interfaces between hot and warm gas with either turbulent mixing (Begelman \\& Fabian 1990; Slavin, Shull, \\& Begelman 1993) or conductive heat transfer occurring in the presence of a magnetic field (Borkowski, Balbus, \\& Fristrom 1990).  This paper is organized as follows. We briefly describe the sample in \\S\\ref{section:Sample} and the observations in \\S\\ref{section:Obs}. We mention the major difficulties encountered in the interpretation of the measurements and discuss the method used to derive column densities in \\S\\ref{section:AnMeth}. We present our findings on the distribution and kinematics of \\ion{O}{6} in \\S\\ref{section:OVIresult}, and the \\ion{O}{6} columns are compared with earlier and contemporary results on other highly-ionized atoms in \\S\\ref{section:ionratios}. We comment on individual sight lines in \\S\\ref{section:peculiar}, discuss our results in \\S\\ref{section:Discussion}, and summarize our conclusions in \\S\\ref{section:Conclusion}. ",
        "conclusions": "}  We have measured the \\ion{O}{6} $\\lambda\\lambda$1031.93, 1037.62 absorption along partial path lengths through the Galactic halo toward 22 Galactic halo stars, and examined the \\ion{O}{6} distribution within $\\sim$3-5 kpc of the Galactic mid-plane. Our results are compared to the findings of Savage {\\it et al}. (2002) and Sembach {\\it et al}. (2002) toward 102 complete sight lines through the Galactic halo. We also study the role of different physical processes in the production of highly-ionized species by comparing the observed N(\\ion{C}{4})/N(\\ion{Si}{4}),  N(\\ion{C}{4})/N(\\ion{O}{6}),  N(\\ion{Si}{4})/N(\\ion{O}{6}), and N(\\ion{N}{5})/N(\\ion{O}{6}) ratios to the predictions of the current models. Our main results are as follows:  1. Strong \\ion{O}{6} absorption is observed at 1031.93 \\AA\\  in all of the directions, except toward HD~97991. We were able to extract the \\ion{O}{6} absorption at 1037.62 \\AA~ for 6 directions. We find that the total \\ion{O}{6} logarithmic column densities vary from 13.65 to 14.57 with an average of $<$log$N$$>$= 14.17$\\pm$0.28 and a median of 14.25. The logarithm of the column density perpendicular to the Galactic mid-plane for these partial paths varies between 13.13 and 14.48 with an average of $<$log$Nsin|b|$$>$= 13.77$\\pm$0.37.  2. The \\ion{O}{6} column densities toward the 22 halo stars are reasonably well described by a patchy exponential distribution. Our measurements are the most consistent with $n_0$= 1.7$\\times$10$^{-8}$ cm$^{-3}$ and a scale height between 2.3 and 4 kpc.  3. We do not see any strong or weak high-velocity component in the \\ion{O}{6} absorption along our sight lines. Sembach {\\it et al}. (2002) reported high velocity \\ion{O}{6} absorption toward $\\sim$60\\% of the complete halo sight lines. The non-detection of high-velocity absorption in our sample suggests that these features originate in the distant halo and confirms the association of the thick disk gas with the low velocity ($|V_{LSR}| \\leq$ 100 \\kms\\ ) absorption in the \\ion{O}{6} survey of Savage {\\it et al}. (2002).  4. Comparison of the \\ion{O}{6} and \\ion{Ar}{1} line centroid velocities reveals a mixed picture. Significant velocity differences ($\\geq$ 15 \\kms\\ ) are observed in 9 cases. \\ion{Ar}{1} and \\ion{O}{6} are closely aligned ($|\\Delta V_{LSR}| \\leq$ 5 km~s$^{-1}$) in 8 cases and there are 5-15 \\kms\\  differences in the remaining 4 cases. The dual behavior of \\ion{O}{6} with respect to \\ion{Ar}{1} may indicate the presence of two types of \\ion{O}{6}-bearing environments along the sight lines studied.  5. There is no obvious correlation between Galactic rotation and the measured LSR velocities of the \\ion{O}{6} absorption. We see deviations from the predicted rotational velocities in positive and negative directions with equal frequency. The \\ion{O}{6}-bearing gas in the low halo participates both in inflow and outflow toward the Galactic plane.  6. The velocity dispersions ($b$-values) of the \\ion{O}{6} profiles vary from 33 to 78 \\kms\\ with an average of 45$\\pm$11 \\kms\\ . These values are much larger than the profile widths $b \\sim$ 18 \\kms\\ expected at T$\\sim$ 3$\\times$10$^5$ K, the temperature at which \\ion{O}{6} peaks in abundance in collisional ionization equilibrium. A considerable amount of turbulent broadening or the presence of multiple components is necessary to explain the observed line breadths. The average \\ion{O}{6} line width is smaller toward the halo stars than it is along complete sight lines through the halo (Savage {\\it et al}. 2002), reflecting the larger number of \\ion{O}{6}-bearing structures along the complete sight lines.  7. Kinematical comparisons of the \\ion{Si}{4}, \\ion{C}{4}, and \\ion{O}{6} line profiles reveal that the \\ion{Si}{4} and \\ion{C}{4} line centroids and line widths correlate with each other well but not with those of the \\ion{O}{6} absorption. There may be a substantial amount of photoionized \\ion{C}{4} and \\ion{Si}{4} without corresponding \\ion{O}{6} along the sight lines.  8. The origin of the highly ionized species in the halo cannot be described by a single physical process. The observed N(\\ion{C}{4})/N(\\ion{Si}{4}), N(\\ion{C}{4})/N(\\ion{O}{6}), N(\\ion{Si}{4})/N(\\ion{O}{6}), and N(\\ion{N}{5})/N(\\ion{O}{6}) ratios along the sight lines suggest that the highly-ionized species are most readily produced by cooling gas in a Galactic fountain flow with contributions from the turbulent interfaces between hot and warm gas. The role of turbulent mixing varies from sight line to sight line.  9. The high N(\\ion{C}{4})/N(\\ion{O}{6}) ratios toward HD~121968 behind the Loop I and IV supernova remnants, and toward HD~177989 behind the Scutum supershell suggest that turbulent mixing has a dominant role in producing \\ion{C}{4} and \\ion{O}{6} toward the disturbed environments of supernova remnants and supershells.  10. Detailed comparisons of the \\ion{C}{4} and \\ion{O}{6} profiles reveal modest to large variations in the N(\\ion{C}{4})/N(\\ion{O}{6}) ratios across the line profiles. We see the largest fluctuations when turbulent mixing has an important role in the production of \\ion{C}{4} and \\ion{O}{6}.  11. The vertical distribution of the high ion ratios toward the halo stars (present analysis) and the complete (Savage {\\it et al}. 2002) sight lines suggest that \\ion{C}{4} and \\ion{Si}{4} trace each other well in the Galactic halo, while \\ion{C}{4} and \\ion{O}{6} exhibit large variations. The average N(\\ion{C}{4})/N(\\ion{O}{6}) ratio is $\\sim$0.6 throughout the entire halo with a decreasing dispersion toward larger $z$ heights.  12. The \\ion{O}{6} absorption does not correlate well with the known \\ion{H}{1} intermediate velocity clouds. We can identify known intermediate velocity features in the \\ion{O}{6} absorption toward HDE~233622, but we see no correspondence toward BD+38~2182 and HD~121800. In the case of HD~121800, the component structure of the \\ion{O}{6} absorption is drastically different from that of the low-ionization states.  13.  Observations toward Galactic stars located in the general direction of $l$= 355$^{\\circ}$ to 360$^{\\circ}$ and $b$= -16$^{\\circ}$ to -21$^{\\circ}$ suggest that the large \\ion{O}{6} columns observed by Savage {\\it et al}. (2002) toward the region $l$= 330$^{\\circ}$ to 36$^{\\circ}$ and $b$= -24$^{\\circ}$ to -33$^{\\circ}$ are related to processes occurring near the Galactic center, and are not associated with the Loop I supernova remnant.  14. Savage {\\it et al}. (2002) found that the general region in the sky toward $l$= 80$^{\\circ}$ to 175$^{\\circ}$ and $b$= -30$^{\\circ}$ to -60$^{\\circ}$ is deficient in \\ion{O}{6}. Our observation toward HD~219188 in this direction reveals an \\ion{O}{6} column similar to those of the nearby extragalactic sight lines. This suggests that the \\ion{O}{6} may be confined to within 2 kpc of the Galactic mid-plane in this direction."
    },
    "0212/astro-ph0212019_arXiv.txt": {
        "abstract": "Over the past few years there has been growing interest in helical magnetic field structures seen at the solar surface, in coronal mass ejections, as well as in the solar wind. Although there is a great deal of randomness in the data, on average the extended structures are mostly left-handed on the northern hemisphere and right-handed on the southern. Surface field structures are also classified as dextral (= right bearing) and sinistral (= left bearing) occurring preferentially in the northern and southern hemispheres respectively. Of particular interest here is a quantitative measurement of the associated emergence rates of helical structures, which translate to magnetic helicity fluxes. In this review, we give a brief survey of what has been found so far and what is expected based on models. Particular emphasis is put on the scale dependence of the associated fields and an attempt is made to estimate the helicity flux of the mean field vs.\\ fluctuating field. ",
        "introduction": "There is now good evidence for the helical nature of the solar magnetic field. Early work by Seehafer (1990) suggested that fitting the line of sight magnetograms of solar active regions to a linear (constant alpha) force-free magnetic field yields systematically negative values of alpha in the northern hemisphere and positive in the southern. Although the evidence for the hemispheric dependence was perhaps not completely convincing back then, subsequent work by different groups (Pevtsov, Canfield, \\& Metcalf 1995; Rust \\& Kumar 1996; Bao et al.\\ 1999; Pevtsov \\& Latushko 2000) have confirmed the initial results.  The quantity being measured in these studies is usually the current helicity, $\\int\\JJ\\cdot\\BB\\,\\dd V$, or the current helicity density, $\\JJ\\cdot\\BB$, where $\\BB$ is the magnetic field strength, $\\JJ=\\nab\\times\\BB/\\mu_0$ is the current density, and $\\mu_0$ is the magnetic permeability. Of particular interest is actually the magnetic helicity $\\int\\AAA\\cdot\\BB\\,\\dd V$, where $\\AAA$ is the magnetic vector potential with $\\BB=\\nab\\times\\AAA$.\\footnote{In general, the boundary of the volume is not a magnetic surface, i.e.\\ $\\BB\\cdot\\nnn\\neq0$, and so $\\int\\AAA\\cdot\\BB\\,\\dd V$ will not be gauge-invariant, i.e.\\ the result will be different if one redefines $\\AAA\\to\\AAA+\\nab\\phi$, where $\\phi$ is an arbitrarily chosen gauge potential. This is why one has instead to use the relative magnetic helicity of Berger \\& Field (1984).}  The magnetic helicity is of great theoretical interest because it satisfies a conservation law: except for small resistive terms, its rate of change depends only on the gains and losses of magnetic helicity through the boundaries. However, the magnetic helicity is a volume integral which is probably hopeless to measure in practice, because the field cannot be observed in the solar interior. What is possible, however, is to measure surface-integrated magnetic helicity fluxes of the form $\\int(\\EE\\times\\AAA)\\cdot\\dd V$, where $\\EE=\\JJ/\\sigma-\\uu\\times\\BB$ is the electric field and $\\sigma$ is the electric conductivity.\\footnote{Again, this quantity is gauge-dependent and has to be substituted by an expression that is compatible with the definition of the relative magnetic helicity of Berger \\& Field (1984).} Regardless of numerous complications, recent work has confirmed the basic hemispheric dependence of the sign of both current helicity densities and surface-integrated magnetic helicity fluxes (Berger \\& Ruzmaikin 2000, DeVore 2000, Chae 2000): negative in the north and positive in the south.  The connection between current helicity and dynamo theory was immediately recognized. R\\\"adler \\& Seehafer (1990) proposed that the observed signs of the current helicity are characteristic of the small scale field rather than the large scale field. From a dynamo point of view this is not only plausible, but also desirable, as we shall explain later. From an observational point of view this is far less obvious, because the field on the scale of active regions and that associated with coronal mass ejections (CMEs) is not generally understood as part of the small scale field. Eclipse images of the sun map out quite clearly the overall field line structure (e.g.\\ Fig.~1 in Low 2001). From these one sees that the field lines in helmet streamers above and around CMEs merges naturally with the large scale of the sun. This is also seen in soft X-ray images from Yohkoh (e.g.\\ Fig.~7 in Low 2001).  The purpose of this paper is to point out that, on theoretical grounds, one might expect a certain degree of simultaneous emergence of small and large scale fields of opposite helicities within each hemisphere. (This concept was also discussed in Blackman \\& Field 2000.) In the following, we explain the reasoning behind such an expectation in light of recent work, and suggest possible observational signatures of the process. ",
        "conclusions": "In this paper we have emphasized the importance of trying to detect simultaneously large and small scale contributions to the losses of helical magnetic fields at the solar surface. The motivation comes mostly from isotropic turbulence simulations (similar high resolution simulations of more realistic settings do not seem to be available yet), but the basic reasoning is sufficiently general to warrant tentative application to the sun. If our picture is correct, it would predict the existence of an as yet unidentified helical component of the magnetic field (with positive magnetic helicity in the northern hemisphere). We expect that this unidentified component should be associated with the large scale field rather than the small scale field. The reason such a component is difficult to detect is related to the fact that the large scale field is not seen directly. It only manifests itself through the systematic orientation of bipolar regions. However, such indirect indications have also been used in the past to estimate the temporal--latitudinal behavior of the large scale magnetic field from synoptic charts (Yoshimura 1976, Stix 1976). This approach should be repeated with more complete recent data to assess at least the sign of the large scale magnetic helicity."
    },
    "0212/astro-ph0212533_arXiv.txt": {
        "abstract": "The sky distribution of cosmic rays with energies above the `GZK cutoff' holds important clues to their origin. The AGASA data, although consistent with isotropy overall, shows evidence for small-angle clustering, and it has been argued that such clusters are aligned with BL Lacertae objects, implicating these as the sources. It has also been suggested that such clusters can arise if the cosmic rays come from the decays of very massive relic particles in the Galactic halo, due to the expected clumping of cold dark matter. We examine these claims and show that both are, in fact, unjustified. ",
        "introduction": "The mystery of the ultra-high energy cosmic rays (UHECRs) with energies exceeding $E_{\\rm GZK}\\simeq4\\times10^{19}$~eV --- the Greisen-Zatsepin-Kuzmin (GZK) `cutoff' \\cite{Greisen:1966jv,Zatsepin:1966jv} --- continues to deepen. This energy sets the threshold for photomeson production on the cosmic microwave background so the observation of such UHECRs (assumed to be protons or heavier nuclei) would indicate that the sources are relatively nearby, within the local supercluster of galaxies \\cite{Stecker:1968uc}. Recent observations by the HiRes air fluorescence detector \\cite{Abu-Zayyad:2002sf} are however inconsistent with previously published data from the Akeno Giant Air Shower Array (AGASA) which ruled out such a cutoff with a significance $\\gtrsim5\\sigma$ \\cite{Takeda:1998ps,Hayashida:2000zr}. HiRes has reported only 1 event above $10^{20}$~eV, whereas about 20 would have been expected on the basis of the AGASA spectrum. The two spectra can be made to agree {\\em below} this energy, if the energies of the AGASA events are systematically lowered by 20\\% (within the quoted uncertainty), however 5 of them still remain above this energy \\cite{Bergman:2002pw}. Subsequently the AGASA collaboration have carefully assessed their energy measurement uncertainties and reaffirmed that their observed spectrum does extend well beyond the GZK energy \\cite{Takeda:2002at}. To resolve this situation requires making simultaneous measurements using both the air shower and air fluorescence methods; such measurements are underway at the Pierre Auger Observatory being constructed in Argentina \\cite{Watson:2001et,pao}.  Another development has been the AGASA observation that the UHECR arrival directions, although consistent with isotropy overall, exhibit clustering on small angular scales \\cite{Takeda:1999sg}. Among the 59 AGASA events above $4\\times10^{19}$~eV, there are 5 `doublets' and 1 `triplet' with separation angle less than the estimated angular resolution of $2.5^\\circ$ \\cite{takeda}.\\footnote{For simulated events with $E>4\\times10^{19}$~eV, 68\\% have a reconstructed arrival direction within $1.8^\\circ$ of the true direction and 90\\% within $3^\\circ$; the corresponding angles for all events above $10^{19}$~eV are $2.8^\\circ$ and $4.6^\\circ$, keeping in mind that the energy resolution is $\\pm30\\%$ \\cite{Hayashida:2000zr,Takeda:2002at}.}  The probability for this to arise by chance from an isotropic distribution is less than 0.1\\%. However this probability is very sensitive to the assumed angular resolution \\cite{Goldberg:2000zq}, e.g. increasing to $\\sim1\\%$ if the angular resolution is $3^\\circ$ \\cite{Anchordoqui:2001qk}. Moreover adding data from three other air shower experiments (Volcano Ranch, Haverah Park, and Yakutsk) {\\em dilutes} the significance. In an earlier such analysis \\cite{Uchihori:1999gu}, 8 doublets and 2 triplets were found in a dataset of 47 AGASA plus 45 other events with $E>4\\times10^{19}$~eV, taking the effective angular resolution of the dataset to be $4^\\circ$. The chance probability for this to arise from an uniform distribution is $\\sim10\\%$, thus statistically not significant.  Nevertheless, the existence of such clusters has been linked to the possibility of (repeating) point sources of UHECR \\cite{Dubovsky:2000gv,Tinyakov:2001ic}, specifically cosmologically distant BL Lacertae \\cite{Tinyakov:2001nr} --- a sub-class of active galactic nuclei (AGN) which have been long discussed as possible accelerators of UHECRs \\cite{Biermann:ep}. However the expected deflections of UHECRs (assumed to be charged particles) by Galactic and intergalactic magnetic fields ought to smear out such tight source correlations \\cite{Stanev:1997,MedinaTanco:1997rt,Alvarez-Muniz:vf}.\\footnote{In fact focussing effects by such fields can give rise to apparent clustering even when the arrival directions are random \\cite{MedinaTanco:1998yq,Harari:2000az,Harari:2002dy}.} Contrary to these results, it has been claimed recently that the correlations with BL Lacs are preserved, even improved, if the UHECRS are protons, after allowing for deflections by the Galactic magnetic field \\cite{Tinyakov:2001ir}. Little is known about the intergalactic magnetic field \\cite{Widrow:2002ud}; requiring rectilinear propagation of protons over the attenuation length of $l\\sim1000$~Mpc at $E>4\\times10^{19}$~eV (decreasing to $\\sim100$~Mpc at $E>10^{20}$~eV \\cite{Stanev:2000fb}) would imply that its homogeneous component on such scales is extremely weak: $B<2\\times10^{-12}(l/1000\\,{\\rm Mpc})^{-1}$~G \\cite{Berezinsky:2002vt}. It has also been claimed \\cite{Blasi:2000ud,Blasi:2001kr} that such clustering is predicted in a model where the UHECR arise from the decay of superheavy relic particles accumulated in the Galactic halo \\cite{Berezinsky:1997hy,Birkel:1998nx}, due to the expected clumping of halo dark matter.  In this paper we examine both these claims in detail, using as our basic statistical tool the two-point correlation function. Our intention is to determine whether the claimed correlations are meaningful, given the present limited event statistics. ",
        "conclusions": "It has been a general expectation that it should be possible to identify the long-sought sources of cosmic rays at energies exceeding $\\sim10^{19}$~eV, when their arrival directions can no longer be randomised by Galactic magnetic fields. However the sky distribution has remained consistent with isotropy up to the highest energies observed. It is clear that at such energies the sources cannot be in the disk of the Galaxy and all experiments indicate that the energy spectrum of such sources is significantly flatter than the component at lower energies. However it is not clear whether the isotropy of arrival directions implicates a relatively {\\em local} population of sources in the Galactic halo (e.g. decaying supermassive dark matter) or a cosmologically {\\em distant} population of astrophysical accelerators (e.g. active galaxies or $\\gamma$-ray bursts). Support for the latter possibility has come from the new HiRes data which shows a cutoff in the spectrum beyond $E_{\\rm GZK}\\simeq4\\times10^{19}$~eV, as has long been expected for extragalactic sources. However the AGASA collaboration has reaffirmed that there is no GZK cutoff in their data, which strongly favours the former possibility.  Given this confusing situation, the indication of small-angle clustering in the AGASA data has naturally been seized upon as a possible further clue as to the nature of the sources. It has been argued both that the clusters coincide with a specific class of extragalactic objects, viz. BL Lacs \\cite{Tinyakov:2001nr,Tinyakov:2001ir}, and, alternatively, that such clustering might arise due to the expected clumping of halo dark matter \\cite{Blasi:2000ud,Blasi:2001kr}. The BL Lac hypthesis is in fact {\\em inconsistent} with the absence of the GZK cutoff in the same AGASA data since many of the identified objects are at very large distances. Hence it appeared more plausible that the sources are clumps of dark matter in the Galactic halo (composed in part of supermassive decaying particles) which would explain the absence of the GZK cutoff.  We have shown that the correlations claimed between BL Lacs and the observed clusters of UHECRs are spurious, being entirely due to selection effects. This is not the first time that such correlations with a particular class of astrophysical accelerators has been claimed; the moral is clearly that care must be taken to not become intrigued by weak accidental correlations and then make arbitrary cuts on the dataset to emphasise them further. We have also found that the extent to which dark matter may be clumped in the halo is not sufficient to generate the observed small-angle clustering, if the UHECRs indeed arise from decaying dark matter. Here the proponents have been misled due to the use of an unphysical density profile for the clumps, as well as a radial distribution in the Galaxy which is inconsistent with the general expectations for hierarchical structure formation.  The net result of our investigations is thus rather negative. The claimed small-angle clustering in the arrival directions of post-GZK UHECRs does not definitively implicate either extragalactic compact sources such as BL Lacs, or decaying clumps of dark matter in the Galactic halo. On the positive side, the forthcoming increase in statistics from the Pierre Auger Observatory will enable us to identify the expected signal from dark matter decays if this is indeed the source of UHECRs. Moreover Auger will also definitively resolve the current contradiction between the air shower and atmospheric fluorescence methods for energy measurement, thus clarifying whether the spectrum does have a GZK cutoff. If so, searches for coincidences with cosmologically distant candidate sources such as active galaxies or $\\gamma$-ray bursts would be of interest. Again the increased statistics provided by Auger would enable the significance of such coincidences to be meaningfully assessed. We will soon know whether the mystery of UHECRs implicates astrophysical sources or new physics beyond the Standard Model."
    },
    "0212/astro-ph0212255_arXiv.txt": {
        "abstract": "The results of astrometric and photometric investigations of the poorly studied open cluster NGC$\\,$1513 are presented.  The proper motions of 333 stars with a root-mean-square error of $1.9\\;{\\rm mas\\,yr^{-1}}$ were obtained by means of  the automated measuring complex \"Fantasy\". Eight astrometric  plates covering the time interval of 101 years were measured and a total of 141 astrometric cluster members identified. $BV$ CCD-photometry was obtained for stars in an area $17\\arcmin\\times 17\\arcmin$ centered on the cluster. Altogether 33 stars with high reliability were considered to be cluster members by two criteria. The estimated age of  NGC\\,1513 is $2.54\\cdot 10^8\\,$ years.\\footnote{Tables~ 2 and 3 are only available in electronic form at the CDS via anonymous ftp to cdsarc.u-strasbg.fr (130.79.128.5) or via http://cdsweb.u-strasbg.fr/cgi-bin/qcat?J/A+A/} ",
        "introduction": "The open cluster NGC\\,1513 is located in the Perseus constellation. Its equatorial and galactic coordinates are:  $\\alpha=4^{\\rm h} 09^{\\rm m} 98^{\\rm s}$,  $\\delta=+49\\degr 31\\arcmin$; $\\ell=152\\fdg 6$, $b=-1\\fdg 57$      $ (2000.0)$.\\\\ In the Tr\\\"umpler(1930)         % catalogue it is classified as II\\,2m a moderately rich cluster with little central concentration. It was investigated astrometrically  by Bronnikova(1958a)           % who  determined the proper motions (PM) using a single pair of plates with an epoch difference of 55 years. Barhatova and Drjahlushina(1960)    % published photographic and photovisual magnitudes of 49 stars from Bronnikova's (1958b)               % list. Del Rio and Huestamendia(1988)  % (RH) obtained photoelectric $UBV$ magnitudes of 31 stars and photographic $RGU$ magnitudes for 116 stars in the cluster region. Spectra and radial velocities of stars in this area have not been investigated. As the epoch difference has increased significantly since the time of the last and only attempt to investigate proper motions of the stars in the region of the open cluster NGC\\,1513 we decided to study it using the material available at present. ",
        "conclusions": "\\setcounter{table}{3} \\begin{table*}  % \\caption{Members and probable cluster members.} \\begin{tabular}{rrccrrccrcc} \\hline  \\multicolumn{5}{c}{Members} & \\multicolumn{6}{c}{Probable members}   \\\\ $N_{pm}$&$N_{phot}$&$M_V$&$(B-V)_0$& $P$& $ N_{phot}$& $M_V$&$(B-V)_0$& $N_{phot}$&$M_V$&$(B-V)_0$ \\\\ 124&  b53&   1.27&  0.09&  59&        4&    2.63&  0.26&     a 69&    3.55&  0.12  \\\\ 126&   37&   1.68&  0.22&  94&       12&    2.30&  0.48&     a 70&    2.79&  0.35  \\\\ 127&   38&   0.85&  0.14&  95&       15&    2.34&  0.31&     a 78&    3.19&  0.34  \\\\ 130&   45&   0.40&  0.09&  83&       16&    1.99&  0.23&     a 80&    3.30&  0.69  \\\\ 132&  b12&   0.33&  0.09&  88&       17&    2.05&  0.12&     a 81&    2.95&  0.48  \\\\ 133&   b7&  -2.18& -0.24&  95&       20&    2.05&  0.36&     a 93&    3.15&  0.58  \\\\ 159&   48&   1.11& -0.02&  90&       23&    2.56&  0.27&     a 97&    3.38&  0.68  \\\\ 160&   76&  -1.15& -0.11&  66&       24&    1.74&  0.30&     a112&    2.56&  0.34  \\\\ 161&   47&   0.64& -0.08&  94&       28&    1.87&  0.23&     a117&    3.47&  0.10  \\\\ 162&   49&   0.19& -0.15&  77&       29&    2.07&  0.30&     a123&    3.51&  0.59  \\\\ 163&   50&  -0.60& -0.07&  94&       33&    2.42&  0.52&     a124&    2.82&  0.26  \\\\ 164&   72&   0.54& -0.06&  93&       34&    2.03&  0.39&     a127&    3.05&  0.37  \\\\ 166&   53&   1.34&  0.05&  93&       35&    2.39&  0.39&     a131&    2.95&  0.16  \\\\ 167&   54&   1.77& -0.01&  80&       39&    2.38&  0.01&     a133&    2.05&  0.37  \\\\ 168&   71&  -1.06& -0.04&  86&       41&    2.53&  0.11&     a157&    3.39&  0.73  \\\\ 169&   55&  -0.64& -0.04&  91&       42&    2.88&  0.39&     a160&    3.29&  0.36  \\\\ 173&   69&   1.60& -0.01&  68&       43&    2.99&  0.40&     a167&    2.66&  0.34  \\\\ 174&   58&   1.50&  0.19&  94&       51&    2.15& -0.07&     a168&    2.97&  0.45  \\\\ 175&   62&  -0.41&  0.88&  93&       52&    2.68&  0.20&     a170&    2.79&  0.30  \\\\ 178&    5&   1.76&  0.18&  84&       56&    2.89&  0.31&     a171&    2.50&  0.30  \\\\ 184&    2&   1.55&  0.05&  66&       57&    2.08&  0.04&     a172&    2.39&  0.20  \\\\ 185&    1&  -0.21& -0.04&  62&       59&    2.16&  0.04&     b  0&    2.83& -0.01  \\\\ 190&    9&   0.77&  0.01&  94&       61&    2.21&  0.16&     b  1&    2.43&  0.25  \\\\ 194&   18&   1.48&  0.21&  91&       70&    1.94&  0.03&     b  2&    2.33&  0.03  \\\\ 215&  110&   0.49&  0.13&  82&       73&    2.55&  0.35&     b  5&    2.83&  0.35  \\\\ 216&  113&   1.51&  0.27&  90&       74&    2.44& -0.07&     b 10&    2.48&  0.49  \\\\ 217&  111&   1.50&  0.18&  89&       75&    2.75&  0.30&     b 13&    1.58&  0.01  \\\\ 219&  106&   1.14&  0.17&  56&       77&    2.69&  0.20&     b 16&    2.36&  0.14  \\\\ 222&   90&   0.66& -0.01&  93&       83&    2.51&  0.39&     b 17&    1.47&  0.15  \\\\ 229&   88&  -1.20&  1.08&  90&       84&    2.65& -0.04&     b 19&    2.99&  0.26  \\\\ 232&   97&   1.82&  0.18&  81&       85&    2.50&  0.07&     b 20&    2.45&  0.42  \\\\ 234&  101&  -0.73& -0.09&  52&       89&    1.52&  0.04&     b 33&    2.51&  0.55  \\\\ 239&   80&   1.69&  0.11&  56&       92&    2.39&  0.23&     b 42&    2.84&  0.37  \\\\ &     &       &      &    &      105&    2.41&  0.32&     b 43&    2.00&  0.39  \\\\ &     &       &      &    &      108&    1.99&  0.24&     b 46&    3.09&  0.42  \\\\ &     &       &      &    &      109&    2.67&  0.18&     b 48&    3.39&  0.14  \\\\ &     &       &      &    &      112&    2.37&  0.10&     b 49&    2.39&  0.17  \\\\ &     &       &      &    &      114&    2.24&  0.25&     b 50&    3.18&  0.21  \\\\ &     &       &      &    &     a  7&    2.02&  0.26&     b 51&    2.24&  0.45  \\\\ &     &       &      &    &     a 12&    2.06&  0.23&     b 56&    3.21&  0.06  \\\\ &     &       &      &    &     a 14&    2.13&  0.24&     b 57&    3.77&  0.50  \\\\ &     &       &      &    &     a 18&    2.11& -0.02&     b 60&    3.41&  0.47  \\\\ &     &       &      &    &     a 21&    3.78&  0.55&     b 65&    3.33&  0.20  \\\\ &     &       &      &    &     a 24&    3.51&  0.16&     b 67&    1.95& -0.03  \\\\ &     &       &      &    &     a 27&    2.63& -0.01&     b 70&    2.50&  0.55  \\\\ &     &       &      &    &     a 28&    3.02&  0.22&     b 74&    3.76&  0.21  \\\\ &     &       &      &    &     a 30&    3.45&  0.09&     b 75&    1.91&  0.25  \\\\ &     &       &      &    &     a 31&    2.88&  0.21&     b 76&    2.45&  0.29  \\\\ &     &       &      &    &     a 39&    3.03&  0.60&     b 79&    3.80&  0.51  \\\\ &     &       &      &    &     a 41&    3.68&  0.13&     b 80&    2.65&  0.48  \\\\ &     &       &      &    &     a 45&    2.63&  0.13&     b 81&    2.06&  0.33  \\\\ &     &       &      &    &     a 53&    2.00&  0.22&     b 89&    2.38&  0.50  \\\\ &     &       &      &    &     a 58&    3.34&  0.54&     b 92&    2.01&  0.22  \\\\ &     &       &      &    &     a 66&    3.55&  0.22&     b101&    2.97&  0.53  \\\\ &     &       &      &    &     a 68&    3.73&  0.10&         &        &        \\\\  \\hline \\end{tabular} \\end{table*}    \\begin{figure} [t] \\centering{ \\vbox{\\psfig{figure=f5.eps,width=8.5cm,height=14cm}}\\par} \\caption []{The $M_V\\&(B-V)_0$ diagramm of all cluster members. ZAMS - Schmidt-Kaler(1965). Isochrone - $z=0.0019, age= 2.54\\cdot 10^8 yr$.} \\end{figure}  All the available $B,V$ - photometry  resulted in the $V$ \\& $B-V$ (CMD) diagram given in Fig.~4. Stars with membership probability $P\\ge 50\\%$ are plotted by solid circles. Obviously it is difficult to define the Main sequence (MS) due to contradictions in the published data on selective absorption in the cluster region (estimates of the mean $E_{B-V}$ vary from 0.52 mag to 0.70 mag). With the absence of U magnitudes of the stars we had to use the ZAMS position from the RH paper. The value of the colour-excess $E_{B-V}$ = 0.67 was calculated from the formula $E_{B-V} = 0.72E_{G-R} -0.005(E_{B-V})^2$ (Steinlin,1968)    % with $E_{G-R}=0.94$ mag and apparent distance modulus $(V-M_V)=12.61$ using the true modulus $(V-M_V)_0 = 10.60$. Stars located on the main sequence within the limits $\\pm 3\\sigma$ are considered to be cluster members according to photometric criterion. In Fig.~5 the  cluster members (according to  photometric and astrometric criteria) are plotted as solid circles. Open circles mark probable cluster members (only by the photometric criterion) as their magnitudes are fainter than the limiting values of the Normal astrograph old plates. In Fig.~5 there are two stars with high membership probabilities at the late evolution stages: N 175: $M_V= -0.41$\\,mag, $(B-V)_0 = 0.88$\\,mag, $P=93\\%$ and N 299 : $M_V=-1.20$\\,mag, $(B-V)_0 = 1.08$\\,mag, $P=90\\%$; ZAMS -  Schmidt-Kaler(1965).                    % The members  and  probable  members are listed in  Table~ 4.  The estimation of the age of the cluster is based on the set of isochrones for stars with mean masses from 0.16 to 7 $M_\\odot$ and different metallicities (Z) published  on the website of the Padove research groupe\\footnote{http://www. pleadi.astro.pd.it} and described in the work of Girardi et al. (2000).  % The superposition of  various isochrones on the CMD shows an almost ideal coincidence of the $2.54\\cdot 10^8$ yr isochrone for $Z=0.019$ with the cluster diagram (Fig.~5). The loop that marks the ending of hydrogen burning in the core, compression  of the core and hydrogen burning in the thick layer is well occupied by stars. The star No 175 is located almost at the base of the red giant branch, and the star No 229 -at the stage of helium burning. Corresponding to this isochrone the turn-off point of the MS is $(B-V)_0 = -0.07$ mag.  It is also interesting to note  that the MS of NGC\\,1513 on the CMD in  absolute coordinates almost completely coinsides with the MS of the nearby cluster NGC\\,1528. Their ages also coincide. The age and MS position of the latter were discussed in the paper by Mermilliod(1981).   % Are these facts accidental? It is still possible that Barhatova(1963)     % who treated them along with NGC\\,1545 cluster as a triplet related by a common origin was right. RH did not confirm the physical connection of NGC\\,1513 and NGC\\,1528 because of the essential difference in their distances. We think that it is worthwhile to obtain data on the membership of stars of NGC\\,1528 not only by the photometric diagram but also by proper motions of its stars. The Pulkovo observatory collection contains a set of negatives taken with the Normal astrograph from 1959 to the present time."
    },
    "0212/astro-ph0212063_arXiv.txt": {
        "abstract": "The linear stability of MHD Taylor-Couette flow of infinite vertical extension   is considered for liquid sodium with its small magnetic Prandtl number Pm of order of $10^{-5}$. The calculations are  performed for a  container with $\\hat\\eta=0.5$ with an axial uniform magnetic field excluding counter rotating cylinders. The sign of the constant $a$ in the basic rotation law  $\\Omega=a+b/R^2$ strongly influences the presented results. It is negative for  resting outer cylinder. The main point here is that   the subcritical excitation which occurs for large Pm disappears for small Pm (cf. Fig. \\ref{minima}).  This is the reason that the existence of the magnetorotational instability remained unknown over decades.  For rotating outer cylinder the limiting case $a=0$ (i.e. $\\hat \\mu = \\hat \\eta^2$) plays an exceptional role. The hydrodynamic instability starts to disappear while the hydromagnetic instability exists with minimal Reynolds numbers  at certain Hartmann numbers of the magnetic field. These  Reynolds numbers exactly scale  with Pm$^{-1/2}$ resulting in moderate values of order $10^4$ for Pm=10$^{-5}$.  However, already for the smallest positive value  of $a$ the Reynolds numbers start to scale as 1/Pm  leading  to  much higher values of order $10^6$ for Pm=10$^{-5}$. Hence, for outer cylinders rotating faster than the limit $a=0$ it is exclusively  the {\\em magnetic}  Reynolds number Rm  which directs the excitation of the instability. They are resulting as lower for insulating walls (`vacuum') than for conducting walls.  Generally, the magnetic Reynolds numbers for liquid sodium have to exceed values of order 10 leading  to frequencies of about   20 Hz for the rotation of the inner cylinder if containers with  (say) 10 cm   radius  are considered. The required magnetic fields are  about 1000 Gauss.   Also  nonaxisymmetric  modes  have  been considered.  With vacuum boundary conditions their excitation is always more difficult than the excitation of   axisymmetric modes; we never observed a crossover of the  lines of marginal stability. For conducting walls, however,  such crossovers exist  for both resting and rotating outer cylinders,  and this  might be essential for future dynamo experiments. In this case, however, the instability  also can onset in form of {\\em oscillating} axisymmetric patterns of flow and field and the Reynolds numbers of these solutions are lower than the Reynolds numbers for the stationary solutions. ",
        "introduction": "The longstanding problem of the generation of turbulence in various hydrodynamically stable situations has found a solution in recent years with the MHD shear flow instability, also called magnetorotational instability (MRI), in which the presence of a magnetic field has a destabilizing effect on a differentially rotating flow with the angular velocity decreasing outwards. The MRI has been  formulated decades ago \\cite{V59,C61} for ideal Taylor-Couette flow, but its importance as the source of turbulence in accretion discs with differential (Keplerian) rotation was first recognized by Balbus and Hawley, \\cite{BH91}.  However, the MRI has never been observed in the laboratory \\cite{DO60,DO62,DC64,B70}. After Goodman and Ji \\cite{GJ01} was the absence of MRI due to the small magnetic Prandtl number approximation used in \\cite{C61}. The magnetic Prandtl number Pm is really very small under laboratory conditions ($\\sim 10^{-5}$ and smaller, see Table \\ref{tab0}).  \\begin{table} \\caption{\\label{tab0} Parameters of the fluids suitable for MHD experiments taken from \\cite{C61} and \\cite{NPCNB02}} \\medskip \\begin{tabular}{|l|c|c|c|c|} \\hline & $\\rho$ [g/cm$^3$] & $\\nu$ [cm$^2$/s] & $\\eta$ [cm$^2$/s] & Pm\\\\[0.5ex] \\hline Mercury & 5.4 & 1.1$\\cdot 10^{-3}$ & 7600 & 1.4$\\cdot 10^{-7}$\\\\[0.5ex] \\hline Gallium & 6.0 & 3.2$\\cdot 10^{-3}$ & 2060 & 1.5$\\cdot 10^{-6}$\\\\[0.5ex] \\hline Sodium & 0.92 & 7.1$\\cdot 10^{-3}$ & 810 & 0.88$\\cdot 10^{-5}$\\\\[0.5ex] \\hline \\end{tabular} \\end{table} A proper understanding of this phenomenon is very important for possible future experiments, including Taylor-Couette flow  dynamo experiments. \\begin{figure} \\psfig{figure=cylinder_new.ps,width=7cm,height=8cm} \\caption{\\label{geometry} Cylinder geometry of the  Taylor-Couette flow with axial magnetic field.} \\end{figure} The simple model of uniform  density fluid contained between two vertically-infinite rotating cylinders is used with constant magnetic field parallel to the rotation axis. For viscous flows the most general form of the rotation law $\\Om(R)$ in the fluid is \\begin{equation} {\\Om}(R) = a+b/{R}^2, \\label{Om} \\end{equation} where $a$ and $b$ are two constants related to the angular velocities $\\Om_{\\textrm{in}}$ and $\\Om_{\\textrm{out}}$ with which the inner and the outer cylinders are rotating and $R$ is the distance from the rotation axis. If $R_{\\textrm{in}}$ and $R_{\\textrm{out}}$ ($R_{\\textrm{out}}>R_{\\textrm{in}}$) are the radii of the two cylinders then \\begin{equation} a={\\hat \\mu-{\\hat\\eta}^2\\over1-{\\hat\\eta}^2}{\\Om}_{\\rm in} \\ \\ \\ \\ \\ {\\textrm{and}} \\  \\ \\ \\ b= R_{\\textrm{in}}^2 {1-\\hat\\mu \\over1-{\\hat\\eta}^2} {\\Om}_{\\rm in} \\label{ab} \\end{equation} with \\begin{equation} \\hat\\mu={\\Om}_{\\textrm{out}}/\\Om_{\\textrm{in}}  \\q  {\\textrm{and}} \\q \\hat\\eta=R_{\\textrm{in}}/R_{\\textrm{out}}. \\label{mu} \\end{equation} Following the Rayleigh stability criterion, $d(R^2 {\\Om})^2/dR>0$, rotation laws are hydrodynamically stable for $a>0$, i.e. $\\hat\\mu>\\hat\\eta^2$. They should in particular  be stable for resting inner cylinder, i.e. $\\hat\\mu \\to \\infty$. Richard and Zahn \\cite{RZ99} focused attention to the experimental results of Wendt \\cite{W1933} who found {\\em nonlinear} instability for this case for Reynolds numbers of order 10$^5$. The finite-amplitude instability of hydrodynamically stable rotation laws must therefore remain in the astrophysical discussion. However, later experiments with very similar Taylor-Couette flow experiments for resting inner cylinder demonstrated  the results of Wendt as due to rather imperfect container constructions and the flow remained laminar even for Reynolds numbers up to 10$^6$,  \\cite{Sch59}.   One of the targets in the present paper is the axisymmetry of the excited modes. We have shown in \\cite{SRS02}  that for containers with conducting boundaries it happens for sufficiently strong magnetic fields that the mode with the lowest eigenvalue (i.e. the lowest Reynolds number) is a nonaxisymmetric mode. As an impressive example,  in Fig. \\ref{old} for Pm=0.01 the crossover of the instability lines for axisymmetric and nonaxisymmetric modes is shown for Hartmann numbers of about 400 (see \\footnote{Note that we considered in \\cite{SRS02}  only  the excitation of the instability for Hartmann number smaller than about 100}). \\begin{figure} \\psfig{figure=p00100-re.ps,width=8cm,height=7cm} \\caption{\\label{old} Instability lines for axisymmetric (solid) and nonaxisymmetric  modes (m=1,  dashed line) for conducting walls and Pm=0.01 ($\\hat\\eta=0.5$).} \\end{figure} Despite of its general meaning this behavior is only known so far for conducting walls and for magnetic Prandtl numbers not smaller than 10$^{-2}$ (see \\cite{KGDIP66}). For possible laboratory experiments we have to extend, however,  the computations to insulating boundaries (vacuum) and to much smaller magnetic Prandtl numbers Pm.   The equations, therefore, are solved here mainly  for the single small magnetic Prandtl number Pm $=10^{-5}$ very close to the value for liquid sodium (see Table \\ref{tab0}). The aspect ratio of the container walls radii in the present paper is fixed to $\\hat\\eta = 0.5$. ",
        "conclusions": ""
    },
    "0212/astro-ph0212549_arXiv.txt": {
        "abstract": "We consider the possibility that an infalling bulge core  or stellar cluster could form an eccentric disk, following tidal disruption by a massive black hole in the center of a galaxy. As a function of central black hole mass, we constrain the core radii and central densities of cluster progenitors capable of becoming nearly Keplerian disks which can support lopsided slow modes. We find that progenitor stellar clusters with core radii less than a pc and densities above a few times $10^5 M_\\odot {\\rm pc}^{-3}$ are likely eccentric disk progenitors near a massive black hole of mass $10^7$ -- $10^8 M_\\odot$. Lower density and larger progenitor cores are capable of causing eccentric stellar disks near more massive black holes. Our constraints on the progenitor cores are consistent with existing N-body simulations, which in one case has produced an eccentric disk. For M31 and NGC 4486B, the estimated progenitor cluster cores are dense and compact compared to Galactic globular clusters, however the cores of nearby galaxies such as M33, M32 and M31 itself are in the right regime. If galaxy mergers can create eccentric disks, then they would be a natural consequence of hierarchical galaxy formation. ",
        "introduction": "Ever since the double-peaked nucleus of M31 was clearly resolved by Hubble Space Telescope (HST) imaging \\citep{lauer}, the morphology of this system has been a challenge to explain.  The most successful kinematic model is the eccentric stellar disk proposed by \\citet{tremaine} where stars in apsidally aligned elliptical orbits about a massive black hole result in an over-density at their apoapses, thus accounting for the brighter peak P1.  The black hole itself resides near the fainter, more centrally located peak denoted P2. Though this model has been successful at matching the observed velocity and luminosity distribution (\\citealt{kormendy,statler,bacon,sambhus}), eccentric disk formation remains a mystery.  \\cite{touma} showed that a small fraction of counter rotating stars, possibly originating from an disrupted globular cluster on a retrograde orbit, could cause a more massive pre-existing stellar disk to develop a lopsided or $m=1$ instability. The self-consistent kinematic modeling of \\citet{sambhus}, which requires a small percentage of counter rotating stars, supports this scenario. \\cite{bacon} proposed that a collision from a passing molecular cloud or globular cluster could knock a pre-existing stellar disk off-center, resulting in a long lived, precessing mode.  The N-body simulations of \\citet{jacobs,taga} have illustrated that lopsided stellar disks could be long-lived.  The scenarios discussed by \\citet{sambhus,touma,jacobs,taga,bacon}, begin with an initially axisymmetric stellar disk, which then becomes lopsided either because of a violent event (such as a collision with a globular cluster) or the growth of an instability.  However \\citet{bekki} proposed that M31's eccentric disk could have been a result of a single disruption event. He succeeded at producing an eccentric stellar disk by disrupting a globular cluster near a massive black hole in an N-body simulation, though an SPH simulation of the disruption of a massive gaseous cloud did not yield an eccentric disk \\citep{bekki2}. The massive black hole in in M31 is moderate with a mass $3 -7 \\times 10^7 M_\\odot$ \\citep{kormendy,bacon}. The eccentric disk itself is nearly as massive as the black hole, $\\sim 3 \\times 10^7 M_\\odot$ \\citep{peng, sambhus, bacon}, presenting a problem for the globular cluster disruption scenario proposed by \\citet{bekki} which assumed that the globular cluster was typical of Galactic globular clusters and of order a million solar masses.    Here we reconsider Bekki's proposal, that a single disruption event could have produced the eccentric disk in M31. The disruption event is likely to be complex so we focus on what final disks are capable of supporting lopsided a slow modes, drawing from the recent work of \\citet{tremaine2001} who developed a formalism for the purpose of predicting the precession rates and eccentricities of discrete modes for low mass, nearly Keplerian disks.  We place limits on the density and core radius of progenitor clusters for eccentric disks near massive black holes. The resulting diagram is useful toward predicting what types of galaxies are most likely to harbor eccentric disks, for predicting the probability that eccentric disk formation events occur and for guiding the initial conditions of N-body simulations which may determine if and how they form. ",
        "conclusions": "In this paper we have considered the possibility that a single disruption event could result in the formation of an eccentric stellar disk such as are found in M31 and possibly NGC 4486B. We explore the scenario of a galaxy core which is stripped as it spirals in toward a nuclear massive black hole until it reaches a critical radius, denoted here as $r_d$, the disruption radius. In the simplest model of an isothermal sphere, the disruption radius is set by the cluster core or King radius and central density.    We suggest that an eccentric disk cannot form unless the core of the galaxy or stellar cluster is disrupted within the sphere of influence of the massive black hole, referred to here as the transition radius, $r_t$, and the core radius is smaller than the radius at which it disrupts. To succeed in making an eccentric disk, we find that the progenitor cluster core must be denser and more compact for less massive black holes.  For the eccentric disk in M31, the progenitor cluster core density must be greater than $10^5 M_\\odot {\\rm pc}^{-3}$ and core radius smaller than a pc. From the radius, width and mass of the most eccentric region of the disk we estimate that the density of the core was a few times $10^6 M_\\odot {\\rm pc}^{-3}$ and the core radius was of order a half pc.  Massive extragalactic globular clusters such as G1 might be dense, compact and massive enough to be an eccentric disk progenitor, however they are probably too diffuse to account for the large disk mass in M31 within a few parsecs of its black hole.   Dense bulges such as found in M31 itself, and in lower luminosity ellipticals galaxies such as M32, are dense and compact enough that they could be similar to progenitors for M31's and NGC 4486B's eccentric disks. Nuclear star clusters, such as found in M33,  may also be dense and compact enough to be progenitors. Though the mass of M33's cluster (a few million $M_\\odot$) is too small to be a progenitor for M31's eccentric disk, the nuclear stellar clusters observed by \\citep{boker2002} in late-type galaxies have half light radii of order 5pc, range in their estimated masses from $10^6-10^8 M_\\odot$, and so could be progenitors for eccentric disks following disruption by a massive black hole. Because of the large mass of bulges and nuclear star clusters at small radii, it would be easier to account for the high masses of the two known eccentric disks by disrupting them, than possible by disrupting a globular cluster. While the massive black hole in M32 might prevent a progenitor similar to M32 from forming an eccentric disk, nuclear clusters such as found in M33 can lack massive black holes and so might provide better progenitor candidates. To date the high angular resolution of HST has resolved extremely high stellar densities or order $10^6 M_\\odot{\\rm pc}^{-3}$ in only the nearest galaxy bulges. Until higher angular resolution observations are available, we will not know if such high density galaxy cores are common.  The two galaxies with double nuclei, M31 and NGC 4486B, exhibit only moderate color variations in their nuclei \\citep{lauer96,lauer98}, a situation that could be a natural consequence of a scenario that involves the merging of galaxy bulges (proposed here), and is more difficult to explain with a scenario that forms a younger disk in situ (an ingredient of the formation scenarios discussed by \\citealt{bacon,touma,taga,jacobs}).  The scenario proposed here is based on a simple tidal disruption argument and can be most quickly tested with N-body simulations such as have been carried out by \\citet{merritt,bekki,holley}. The simulations of \\citet{bekki} established that the disruption of a cluster could result in the formation of an eccentric disk, and \\citet{merritt,holley} have carried out simulations based on realistic galaxy profiles and with massive black holes.    \\citet{holley} established that a spinning nuclear stellar disk can be a remnant. It remains to be seen whether the merger of two galaxy cores can result in the creation of an eccentric stellar disk.  We suspect that the merger of a primary galaxy containing a very massive black hole with a secondary with a `core-type' or shallow central surface brightness profile, and significantly lower mass black hole, would be most likely to form an eccentric disk with N-body simulations. For the core-type galaxies, the break radius and density at this radius provide an equivalent for the King or core radius and central density used in Figure 1 and so can be used to estimate the likelihood that the core can disrupt to form an eccentric disk.   If the merger of galaxy bulges can result in the formation of an eccentric disk, then their formation would be a natural consequence of hierarchical galaxy formation, and can also be used to probe the properties of the parents of galaxies which contain them."
    },
    "0212/astro-ph0212313_arXiv.txt": {
        "abstract": "The launch of high-spectral-resolution x-ray telescopes (\\textit{Chandra}, \\textit{XMM}) has provided a host of new spectral line diagnostics for the astrophysics community.  In this paper we discuss Doppler-broadened emission line profiles from highly supersonic outflows of massive stars.  These outflows, or winds, are driven by radiation pressure and carry a tremendous amount of kinetic energy, which can be converted to x rays by shock-heating even a small fraction of the wind plasma.  The unshocked, cold wind is a source of continuum opacity to the x rays generated in the shock-heated portion of the wind.  Thus the emergent line profiles are affected by transport through a two-component, moving, optically thick medium. While complicated, the interactions among these physical effects can provide quantitative information about the spatial distribution and velocity of the x-ray-emitting and absorbing plasma in stellar winds. We present quantitative models of both a spherically-symmetric wind and a wind with hot plasma confined in an equatorial disk by a dipole magnetic field. ",
        "introduction": "Ultraviolet spectra of O and B stars (with luminosities up to $L=10^6 \\;\\mathrm{L_{\\odot}}$ and surface temperatures $T \\gtrsim 3 \\;\\mathrm{T_{\\odot}}$) show the signatures of rapidly expanding winds, with velocities on the order of a few $1000\\; \\mathrm{km}\\; \\mathrm{s}^{-1}$, densities of order $10^{10}\\; \\mathrm{cm}^{-3}$, and mass-loss rates up to $10^{-5}\\;\\mathrm{M_{\\odot} \\;\\mathrm{yr}}^{-1}$. These stars are detected to have soft-x-ray luminosities of $~10^{-7}$ times their total (bolometric) luminosities.\\cite{1989ApJ...341..427C} In cooler stars, like the Sun, x rays are produced in coronae, which are high-temperature regions near the surface of the star, and their x-ray line widths are primarily indicative of thermal broadening. The line widths observed in x-ray spectra of hot stars, however, provide evidence that the emission originates in extended, high-velocity winds. \\cite{2001A&A...365L.312K}  The \\textit{Chandra X-ray Observatory} has the spectral resolution to resolve Doppler-broadened lines from hot stars, providing quantitative information about the geometry and kinematics of the wind and putting important constraints on physical models that may explain the observed x rays. \\textit{Chandra} uses a set of concentric parabolic and hyperbolic mirrors with an effective area of about $10\\; \\mathrm{cm}^2$ to focus x rays on its instrument package. The Medium Energy Grating (MEG) provides a resolution of $\\Delta \\lambda = 0.023\\;\\mathrm{\\mbox{\\AA}}$ from $0.4$ to $5.0\\;\\mathrm{keV}$ ($31$ to $2.5\\;\\mathrm{\\mbox{\\AA}}$). ",
        "conclusions": ""
    },
    "0212/astro-ph0212131_arXiv.txt": {
        "abstract": "{Pure N-body high-resolution re-simulations of 13 massive dark matter halos in a $\\Lambda$CDM cosmology are presented. The resulting sample is used to investigate the physical origin of the Fundamental Plane, ``Faber-Jackson'', and ``Kormendy'' relations observed for nearby galaxy clusters. In particular, we focus on the role of dissipationless hierarchical merging in establishing or modifying these relations. Contrarily to what found in the case of the Faber-Jackson and Kormendy relations for {\\it galaxies} (see Londrillo, Nipoti \\& Ciotti, this conference), dissipationless merging on {\\it cluster scales} produces scaling relations remarkably similar to the observed ones. This suggests that gas dissipation plays a minor role in the formation of galaxy clusters, and the hierarchical merger scenario is not at odd with the observed regularity of these systems. ",
        "introduction": "Like early-type galaxies, also galaxy clusters follow well defined scaling relations among their main observables: in particular, a luminosity--velocity dispersion relation (i.e., a ``Faber-Jackson--like'' relation), a luminosity--radius relation (i.e., a ``Kormendy--like'' relation), and the  Fundamental Plane (hereafter FP; e.g., Schaeffer et al. 1993, hereafter S93; Adami et al. 1998; and for ellipticals: Djorgovski et al. 1987; Dressler et al. 1987; Faber \\& Jackson 1976; Kormendy 1977). Besides their potential importance as distance indicators, these relations contain information on the cluster formation and evolution processes. Within the commonly accepted cosmological scenario, where cold dark matter (CDM) is the dominant component of the Universe, structure formation is driven by hierarchical dissipationless merging: small DM halos form first, then they merge together to form larger systems, and the evolution of the baryonic matter follows that of the hosting DM halos.  The importance of investigating the physical origin of the observed scaling relations, and the role of dissipationless hierarchical merging in establishing or modifying them, is therefore apparent, and explains the large number of theoretical works devoted to this problem in the case of elliptical galaxies (e.g. Capelato et al. 1995; Nipoti, Londrillo, Ciotti 2002a,b; Gonzales \\& van Albada 2002; Evstigneeva et al. 2002; Dantas et al. 2002; Londrillo et al., this conference).  In particular, it has been shown that, while the FP is reproduced by the end--products of equal mass mergers, the Faber--Jackson and the Kormendy relations are {\\em not} (Nipoti et al. 2002ab; Londrillo et al., this conference).  In the case of galaxy clusters however, almost no theoretical studies are devoted to their FP (see, e.g., Pentericci, Ciotti \\& Renzini 1995; Fujita \\& Takahara 1999; Beisbart, Valdarnini \\& Buchert 2001).  We address this problem in the present work, by means of high resolution re-simulations of very massive DM halos (the hosts of galaxy clusters) in a $\\Lambda$CDM cosmology.  Following their hierarchical dissipationless evolution in detail, we verify whether they reproduce the observed scaling relations, both at the present time, and at higher redshift.  A full analysis of this work will be given elsewhere (Lanzoni et al., in preparation), while we present here a selection of the main astrophysical results, focusing on the numerical issues. ",
        "conclusions": "We have described a technique for increasing the (mass and spatial) resolution of objects selected in a given cosmological simulation. Applying it to sample of 13 massive halos at $z=0$ in a $\\Lambda$CDM cosmology, we have investigated whether the scaling relations observed for nearby galaxy clusters are also followed by their DM hosts, in a dissipationless hierarchical merging scenario.  By considering the most massive progenitors at $z=0.5$ and $z=1$ of the 13 selected halos, we have also investigated whether the same scaling relations were already in place at those earlier epochs.  The main conclusions can be summarized as follows: \\begin{itemize}  \\item the high-resolution re-simulation technique is very powerful and reliable. The precision of the method depends on the number of particles composing the original object, and on the dimension of the boundary region considered to avoid the ``contamination'' from macroparticles.  \\item The DM hosts of galaxy clusters do define a FP that is practically indistinguishable from the one observed for nearby galaxy clusters.  \\item A very good agreement between the scaling relations of simulated and observed clusters is also found in terms of the $L$-$\\sigma$ and the $L$-$R$ relations.  \\item The FP of simulated clusters is already defined at $z=0.5$ and $z=1$, and it is very similar and as thin as the one observed locally; the only differences are due to an obvious shift towards smaller radii and masses for increasing $z$. Also the $L$-$\\sigma$ and $L$-$R$ relations are already in place at high redshifts, with a difference in the slope (with respect to the relations observed locally) due to the fact that at recent epochs, more massive objects evolve more rapidly than smaller systems.  \\end{itemize}  All the results about the scaling relations of simulated and observed clusters are obtained by assuming a mass-to-light ratio with mean value of the order of $450 h \\msol/\\lsol$ (well in agreement with what estimated from observations of galaxy clusters), and with a dependence on the luminosity also suggested by the observations.  Given the very different nature of the two components (baryonic and non baryonic) and of the physical processes acting on them, the fact that the DM halos define the same relations shown by real clusters is not obvious at all.  Moreover, dissipationless hierarchical merging at galactic scales is unable to reproduce the observed scaling relations of elliptical galaxies: the end-products of major mergers do lie on the observed FP, but substantial deviations from the Faber-Jackson and the Kormendy relations are found in this case (see Londrillo et al., this conference).  Such a result suggests that, while dissipation has a major role in the formation and evolution of ellipticals, it is negligible with respect to the pure gravity in settling the properties of galaxy clusters."
    },
    "0212/astro-ph0212307_arXiv.txt": {
        "abstract": "We present subarcsecond resolution mid infrared images of NGC 4151 at 10.8$% \\micron$ and 18.2$\\micron$. These images were taken with the University of Florida mid-IR camera/spectrometer OSCIR at the Gemini North 8 m telescope. We resolve emission at both 10.8$\\micron$ and 18.2$\\micron$ extending $% \\thicksim 3.5\\arcsec$ across at a P.A. of $\\thicksim 60\\degr$. This coincides with the the narrow line region\\ of NGC 4151 as observed in [OIII] by the Hubble Space Telescope. The most likely explanation for this extended mid-IR emission is dust in the narrow line region heated by a central engine. We find no extended emission associated with the proposed torus and place an upper limit on its mid-IR size of $\\lesssim 35$ pc. ",
        "introduction": "NGC 4151 is one of the nearest (13.2 Mpc, H$_{o}=75$ km s$^{-1}$ Mpc$^{-1}$) and best studied active galactic nuclei (AGN). The nucleus hosts a highly variable continuum and line emission source. Continuum variability, first reported by Fitch et al. (1967), has been observed at several wavelengths including X-ray (Papadakis et al. 1995), UV (Clavel et al. 1990), and optical (Lyutyi 1972). Classified as a Seyfert 1.5 by Osterbrock \\& Koski (1976), NGC\\ 4151 displayed characteristics of a Seyfert 2 (Penston \\& P{$% \\acute{e}$}rez 1984) during a low luminosity state in 1984, and at a later date characteristics of a Seyfert 1 (Ayani \\& Maehara 1991).  The mid-infrared emission from NGC\\ 4151 has been suggested to arise from either thermal emission from dust grains or synchrotron emission. Discussions of the thermal vs. nonthermal origin of the infrared emission in NGC 4151 can be found in Rieke \\& Lebofsky (1981), Edelson \\& Malkan (1986), Carelton et al. (1987), Edelson et al. (1987), and de Kool \\& Begelman (1989). A direct method to investigate the origin of the mid-IR emission mechanism, as proposed by Neugebauer et al. (1990) (hereafter N90), is a measurement of the size of the emitting region. They suggest that a nonthermal self-absorbed synchrotron source would be $<1$mas, and hence unresolvable. However, if the mid-IR emission is due to heated dust grains, the size of the region would be $>0.1\\arcsec$.  Observations show that the mid-IR emission in NGC 4151 is compact. Comparison between 60$\\arcsec$ resolution IRAS 12 $\\micron$ flux density measurements and 6$\\arcsec$ resolution ground-based 10.6 $\\micron$ measurements agree to within $\\thicksim 6\\%$ (Edelson et al. 1987). Mid-IR observations by Rieke \\& Low 1972; Rieke \\& Lebofsky 1981; Ward et al. 1987 also did not detect any extended emission with resolutions $\\gtrsim 6\\arcsec$% . Observations by ISO (Infrared Space Observatory; Rodriquez-Espinosa et al. 1996 - hereafter RE96) show a strong warm dust component in NGC\\ 4151 and suggest a thermal origin from a geometrically and optically thick dusty torus and/or a dusty narrow line region (NLR). These observations however were at a resolution of $180\\arcsec$. Using a technique of near-simultaneous north-south scans at 2.2 $\\micron$ and 11.2 $\\micron$, N90 was able to resolve the 11.2 $\\micron$ emitting region to be 0\\farcs16$\\pm 0$\\farcs04. However this technique measured the size in only one spatial direction and was insufficient to explore the mid-IR morphology of the circumnuclear region. In addition, these north-south scans could not investigate the NLR of NGC 4151 which is primarily orientated in an east-west direction. In this paper we present high resolution mid-IR imaging which, to the best of our knowledge, resolves the inner NLR\\ of NGC 4151 for the first time at 10 $% \\micron$ and 18 $\\micron$. ",
        "conclusions": "In this paper we have used mid-IR\\ imaging and first order radiative transfer analysis assuming dust grains in thermal equilibrium to study the central $\\thicksim 11\\arcsec$ of NGC 4151. Our conclusion are as follows.  \\  1. We detect extended mid-IR emission at 10.8 $\\micron$ and 18.2 $\\micron$ in the circumnuclear region of NGC 4151. This emission extends approximately $\\thicksim 200$ pc ($\\thicksim $ $3.5\\arcsec$) at a P.A. $\\thicksim 60\\degr$ correlating with the NLR\\ region as seen in [OIII] $\\lambda $5007 by Evans et al. (1993) and Kaiser et al. (2000) using HST.  2. With the PSF\\ scaled to 100\\% of the peak of NGC 4151 we measure limits to the unresolved and resolved components of $\\leq 73\\%$ and $\\geq 27\\%$ respectively.  3. Mid-IR line emission contributes $<10\\%$ of the extended emission at 10.8 $\\micron$ and 18.2 $\\micron$. The lack of any detectable PAH\\ emission also shows that star formation is weak in the circumnuclear region.  4. Assuming that the luminosity in NGC 4151 is anisotropic ($\\thicksim 13\\times $), the extended mid-IR emission in NGC 4151 is consistent with thermal re-radiation from dust grains in the NLR heated by a central engine.  5. We place an upper limit on the size of the torus in the mid-IR of $% \\lesssim 35$pc consistent with the measurements of N90, and Ruiz et. al. (2002). This results in an upper limit to the mid-IR contribution from a dusty torus in NGC 4151 of $\\leq 73\\%$ of the total emission at 10.8 $% \\micron $ and 18.2 $\\micron$ based on our unresolved (PSF) component.  6. Mid-IR measurements of the proposed torus by N90 as well as upper limits derived from this paper are roughly consistent with the ``onion-skin'' model of Pedlar et al.\\ (1998). In this model, ionizing photons in the plane of the torus may be $\\thicksim $10 - 40 times less than seen from Earth."
    },
    "0212/astro-ph0212077_arXiv.txt": {
        "abstract": "We construct hydrogen atmosphere models for magnetized neutron stars in radiative equilibrium with surface fields $B=10^{12}-5\\times 10^{14}$~G and effective temperatures $\\Teff\\sim\\mbox{a few}\\times 10^5-10^6$~K by solving the full radiative transfer equations for both polarization modes in the magnetized hydrogen plasma. The atmospheres directly determine the characteristics of thermal emission from isolated neutron stars. We study the effects of vacuum polarization and bound atoms on the atmosphere structure and spectra. For the lower magnetic field models ($B\\sim 10^{12}$~G), the spectral features due to neutral atoms lie at extreme UV and very soft X-ray energies and therefore are not likely to be observed.  However, the continuum flux is also different from the fully ionized case, especially at lower energies. For the higher magnetic field models, we find that vacuum polarization softens the high energy tail of the thermal spectrum.  We show that this depression of continuum flux strongly suppresses not only the proton cyclotron line but also spectral features due to bound species; therefore spectral lines or features in thermal radiation are more difficult to observe when the neutron star magnetic field is $\\go 10^{14}$~G. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/astro-ph0212241_arXiv.txt": {
        "abstract": "In quintessence models, the dark energy content of the universe is described by a slowly rolling scalar field whose pressure and energy density obey an equation of state of the form $p=w \\rho$;  $w$ is in general a function of time such that $w<-1/3$, in order to drive the observed acceleration of the Universe today. The cosmological constant model ($\\Lambda$CDM) corresponds to the limiting case $w=-1$. In this paper, we explore the prospects of using the Lyman-$\\alpha$ forest to constrain $w$, using semi-analytical techniques to model the intergalactic medium (IGM).  A different value of $w$ changes both the growth factor and the Hubble parameter as a function of time. The resulting change in the optical depth distribution affects the optical depth power spectrum, the number of regions of high transmission per unit redshift and the cross-correlation coefficient of spectra of quasar pairs. These can be detected in current data, provided we have independent estimates of the thermal state of the IGM, its ionization parameter and the baryon density. ",
        "introduction": "\\label{intro} Observed cosmic microwave background anisotropies convincingly demonstrate that the universe is spatially flat (de Bernardis \\et 2002; Netterfield \\et 2002; Pryke \\et 2002). The luminosity-distance relation, as determined from high redshift type Ia Supernovae (Garnavich \\et 1998; Riess \\et 1998; Perlmutter \\et 1999), requires a spatially flat universe to be currently dominated by some type of dark energy, which is nearly homogeneous and has negative pressure, such as for example a cosmological constant. Several independent lines of argument also favour a low-density, vacuum-energy dominated universe, for example the abundance of clusters (Bahcall \\et 2003) and their X-ray properties (Allen \\et 2002), the clustering of galaxies (Efstathiou \\et 2002; Verde \\et 2002), estimates of the age of the universe and the current value of the Hubble parameter.  The energy density associated with the cosmological constant may actually decay in time. Such quintessence models (e.g. Caldwell \\et 1998; henceforth QCDM models) have an equation of state $p=w\\, \\rho$, with $w<-1/3$, where $w=-1$ corresponds to the more familiar cosmological constant model. Recently, a number of observational tests have been proposed to measure $w$ and its redshift dependence (Kujat \\et 2001; Wang \\& Garnavich 2001; Matsubara \\& Szalay 2002; Gerke \\& Efstathiou 2002).  In this {\\it Letter} we focus on the prospects of using the \\lya forest to constrain $w$. The \\lya forest region in quasar (QSO) spectra contains hundreds of hydrogen \\lya absorption lines (see Rauch 1998 for a review), most of which are produced by small density fluctuations in the intervening intergalactic medium (Cen \\et 1994; Bi \\& Davidsen 1997, hereafter BD97). Since these structures are still reasonably linear, they are good probes of the underlying large-scale matter distribution, and hence are sensitive to the cosmological model. For example, Hui \\et (1999) and McDonald (2001) suggested to apply the Alcock-Paczinski test (Alcock \\& Paczinski 1979) to estimate the cosmological constant and Croft \\et (2002) constrained the dark matter power spectrum on large scales. The \\lya forest is promising in that it can be used over a larger redshift range than most other potential probes of $w$.  Here we use semi-analytical models of the \\lya forest, which have been shown to be successful in reproducing reasonably well most of its observed properties, such as the column density distribution function, the line-widths distribution and the number of lines per unit redshift (BD97; Viel \\et 2002a, 2002b; Roy Choudury \\et 2001).  Hydrodynamical simulations are required for more accurate predictions on smaller scales (e.g. Theuns \\et 1998, 2002; Dav\\'e \\et 1999; Bryan \\et 1999). We examine several statistics of the flux spectrum, such as the probability distribution function of the optical depth, the optical depth power spectrum, the number of underdense regions per unit redshift and the cross-correlation coefficient of spectra of quasar pairs. The first two statistics are not directly observable but they can be recovered from the flux distribution by looking at higher order lines through pixel optical depths techniques (Sect. 3). Section~2 briefly reviews the properties of quintessence models and describes our simulation techniques. Section~3 contains our results and a brief discussion. ",
        "conclusions": ""
    },
    "0212/astro-ph0212288_arXiv.txt": {
        "abstract": "Interstellar magnetic fields are strong: up to 25$\\mu$G in spiral arms and 40$\\mu$G in nuclear regions. In the spiral galaxy NGC~6946 the average magnetic energy density exceeds that of the thermal gas. Magnetic fields control the evolution of dense clouds and possibly the global star formation efficiency in galaxies. Gas flows and shocks in spiral arms and bars are modified by magnetic fields. Magnetic forces in star-forming circumnuclear regions are able to drive mass inflow towards the active nucleus. Magnetic fields are essential for the propagation of cosmic rays and the formation of galactic winds and halos. ",
        "introduction": "{\\it ... but it's easier than most people think.} In radio continuum the typical degrees of polarization are much higher than in the other spectral ranges, and we benefit from the development of large instruments and sensitive receivers. This is why most of our knowledge on interstellar magnetic fields in our Galaxy (\\opencite{Reich1994}; \\opencite{Heiles1996}; \\opencite{Beck2001}) and in external galaxies is based on polarized radio emission and its Faraday rotation (see also \\opencite{Beck2000}; \\opencite{Beck2002}). To detect extended features and achieve high resolution, data from interferometric (synthesis) and single dish telescopes have to be combined. ",
        "conclusions": ""
    },
    "0212/hep-ph0212404_arXiv.txt": {
        "abstract": "We solve the nonequilibrium dynamics of a $3+1$ dimensional theory with Dirac fermions coupled to scalars via a chirally invariant Yukawa interaction. The results are obtained from a systematic coupling expansion of the 2PI effective action to lowest non-trivial order, which includes scattering as well as memory and off-shell effects. The dynamics is solved numerically without further approximation, for different far-from-equilibrium initial conditions. The late-time behavior is demonstrated to be insensitive to the details of the initial conditions and to be uniquely determined by the initial energy density. Moreover, we show that at late time the system is very well characterized by a thermal ensemble. In particular, we are able to observe the emergence of Fermi--Dirac and Bose--Einstein distributions from the nonequilibrium dynamics. ",
        "introduction": "\\label{soverview}  The abundance of experimental data on matter in extreme conditions from relativistic heavy-ion collision experiments, as well as applications in astrophysics and cosmology has lead to a strong increase of interest in the dynamics of quantum fields out of equilibrium. Experimental indications of thermalization in collision experiments and the justification of current predictions based on equilibrium thermodynamics, local equilibrium or \\mbox{(non--)linear} response pose major open questions for our theoretical understanding~\\cite{Braun-Munzinger:2001mh,Serreau:2002yr}. One way to resolve these questions is to try to understand quantitatively the far-from-equilibrium dynamics of quantum fields, without relying on the assumption of small departures from equilibrium, or on a possible separation of scales that forms the basis of effective kinetic descriptions \\cite{Blaizot:2001nr}. In contrast to close-to-equilibrium approaches, the quantum-statistical fluctuations of the fields are not assumed to be described by a thermal ensemble. Moreover, one goes beyond the range of applicability of the usual gradient expansion and dilute-gas approximation.   In contrast to thermal equilibrium, which keeps no information about the past, nonequilibrium dynamics poses an initial value problem: time-translation invariance is explicitly broken by the presence of the initial time, where the system has been prepared. The question of thermalization investigates how the system effectively looses the dependence on the details of the initial condition, and becomes approximately time-translation invariant at late times. According to basic principles of equilibrium thermodynamics, the thermal solution is universal in the sense that it is independent of the details of the initial condition and is uniquely determined by the values of the (conserved) energy density and of possible conserved charges.\\footnote{Here we consider closed systems without coupling to a heat bath or external fields, which could provide sources or sinks of energy.} There are two distinct classes of universal behavior, corresponding to Bose-Einstein and Fermi-Dirac statistics respectively.  The description of the effective loss of initial conditions and subsequent approach to thermal equilibrium in quantum field theory requires calculations beyond so--called ``Gaussian'' (leading-order large--$N$, Hartree or mean-field type) approximations \\cite{gaussian1,gaussian2a,gaussian2b}. Similar to the free-field theory limit, these approximations typically exhibit an infinite number of additional conserved quantities which are not present in the underlying interacting theory \\cite{Cooper:1996ii,Aarts:2001wi}.  These spurious constants of motion constrain the time evolution and lead to a non-universal late-time behavior \\cite{Berges:2001fi}. It has recently been demonstrated in the context of scalar field theories, that the approach to quantum thermal equilibrium can be described by going beyond these approximations \\cite{Berges:2000ur,Berges:2001fi}. In particular, this has been achieved by using a systematic coupling expansion \\cite{Berges:2000ur} or a $1/N$ expansion to next-to-leading-order \\cite{Berges:2001fi,Aarts:2002dj} of the two-particle irreducible generating functional for Green's functions, the so-called 2PI effective action \\cite{Baym,Cornwall:1974vz,Calzetta:1988cq,Ivanov:1998nv}.  In this context, the corresponding equations of motion have been shown to lead to a universal late time behavior, in the sense mentioned above, without to have recourse to any kind of coarse-graining.  In this work, we study the far-from-equilibrium time evolution of relativistic fermionic fields and their subsequent approach to thermal equilibrium. Nonequilibrium behavior of fermionic fields has been previously addressed in several scenarios \\cite{gaussian2a,gaussian2b,Danielewicz:kk,Greene:2000ew,Joyce:2000ed,Lawrie:2000jg}, including the full dynamical problem with fermions coupled to inhomogeneous {\\em classical$\\,$} bosonic fields \\cite{Aarts:1998td}.  Here we go beyond these approximations and compute the nonequilibrium evolution in a $3+1$ dimensional {\\em quantum} field theory of Dirac fermions coupled to scalars in a chirally invariant way.  As a consequence, we are able to study the approach to quantum thermal equilibrium. For the considered nonequilibrium initial conditions we find that the late-time behavior is universal and characterized by Fermi--Dirac and Bose--Einstein distributions, respectively.  The results are obtained from a systematic coupling expansion of the 2PI effective action to lowest nontrivial (two-loop) order, which includes scattering as well as memory and off-shell effects.  The nonequilibrium dynamics is solved numerically without further approximations. We emphasize that, given the limitations of a weak-coupling expansion, this is a first-principle calculation with no other input than the dynamics dictated by the considered quantum field theory for given initial conditions.  In Sect.~\\ref{Sect2PIeffaction}, we review the 2PI effective action for fermions, which we use to derive exact time evolution equations for the spectral function and the statistical two-point function in Sect.~\\ref{SectevolrhoF}. We discuss the Lorentz structure of our equations in Sect.~\\ref{sec:lorentz} and exploit some symmetries in Sect.~\\ref{sec:symmetries}. In Sect.~\\ref{sec:model}, we specify to a chiral quark model for which we solve the nonequilibrium dynamics. The initial conditions are discussed in Sect.~\\ref{sec:initial} and some details concerning the numerical implementation are given in Sect.~\\ref{sec:numerics}. The numerical results are presented and discussed in Sect.~\\ref{Sectfarfromequilibrium} and Sect.~\\ref{sec:statistics}. We attach two appendices discussing in detail the quasiparticle picture we used to interprete some of our results. ",
        "conclusions": "Conclusions}    In this paper we have discussed the far-from-equilibrium dynamics and subsequent thermalization of a system of coupled fermionic and bosonic quantum fields. We solved the nonequilibrium dynamics beyond mean-field type approximations by calculating the complete lowest non-trivial order in a systematic coupling expansion of the 2PI effective action, which includes direct scattering as well as memory and off-shell effects. To our knowledge this is the first time that such a calculation is performed without further approximations. As a result, we show that, for various far-from-equilibrium initial conditions, the late-time behavior is universal and uniquely determined by the value of the initial energy density. Moreover, we are able to probe the approach to quantum thermal equilibrium, characterized by the emergence of Fermi--Dirac and Bose--Einstein distribution functions, with high accuracy. We emphasize that in the present calculation, besides the limitations of a coupling expansion, there is no other input than the dynamics dictated by the considered quantum field theory for given nonequilibrium initial conditions.  This work can be extended in many directions. Most of the equations we derive are also valid in the phase with spontaneous symmetry breaking. Combined with earlier work on scalar theories \\cite{Berges:2001fi,Aarts:2002dj,Berges:2002cz}, this provides a description of the nonequilibrium dynamics of the linear $\\sigma$-model for QCD with two quark flavors. The model has served for many years as a valuable testing ground for ideas on the equilibrium phase structure of low energy QCD at nonzero temperature and density. With the present techniques a quantitative understanding of the out-of-equilibrium physics of this model is within reach.       \\subsection*{Acknowledgment} We thank C.~Wetterich for many discussions and collaboration on related work. We are also grateful to W.~Wetzel for his continuous support with computers. Sz.~B.\\ acknowledges the hospitality of the Institute f\\\"ur Theoretische Physik, Heidelberg. His work was supported by the short-term scholarship program of the DAAD.   \\appendix"
    },
    "0212/astro-ph0212527_arXiv.txt": {
        "abstract": "We present theoretical UBVI- and bolometric light curves of SNe Ia for several explosion models, computed with our multi-group radiation hydro code. We employ our new corrected treatment for line opacity in the expanding medium. The results are compared with observed light curves. Our goal is to find the most viable thermonuclear SN model that gives good fits not only to a typical SN Ia light curve, but also to X-ray observations of young SNIa remnants. It appears that classical 1D SNIa models, such as deflagration W7 and delayed detonation DD4, fit the light curves not so good as a new 3D deflagration model by Reinecke et al (which is averaged over angles for our LC modelling). This model seems good also in reproducing X-ray observations of Tycho SNR. We believe that the main feature of this model which allows us to get correct radiation during the first month, as well as after a few hundred years, when an SNR forms, is strong mixing that pushes material enriched in iron and nickel to the outermost layers of SN ejecta. ",
        "introduction": "At the moment, there are many models of thermonuclear explosion of a star, that lead to the event we know as a Type~Ia Supernova (SN~Ia). Some of them were discussed in the talk by J.Niemeyer~\\cite{sorokina.Nhere}. Only a few parameters, such as kinetic energy and total $^{56}$Ni production, can be derived directly from the explosion modelling and compared with the observational values. The subsequent evolution of the exploded star gives us much more possibilities to compare models and to decide which one fits observations better by reproducing more details in SN~Ia light curves and spectra. We will focus here on the broad-band UBVI and bolometric light curve computations for SN~Ia models.  There are several effects in SNe physics which lead to difficulties in the light curve modelling of any type of SNe. For instance, an account should be taken correctly for deposition of gamma photons produced in decays of radioactive isotopes, mostly $^{56}$Ni and  $^{56}$Co. After being emitted, gamma photons travel through the ejecta and can finish up in either thermalization or in non-coherent scattering processes. To find this one has to solve the transfer equation for gamma photons together with hydrodynamical equations. Full system of equations should involve also radiative transfer equations in the range from soft X-rays to infrared for the expanding medium. There are millions of spectral lines that form SN spectra, and it is not a trivial problem to find a convenient way how to treat them even in the static case. The expansion makes the problem much more difficult to solve: hundreds or even thousands of lines give their input into emission and absorption at each frequency.  On the first glance, modelling of SNe~Ia seems easier than of other types of supernovae, since the hydro part is very simple: coasting stage starts very early, there are no shocks, and no additional heating from them.  On the other hand, much more difficulties arise in the radiation part. SNe~Ia becomes almost transparent in continuum at the age of a few weeks. This means that NLTE effects are stronger than for other types of supernovae. Radiation is decoupled from matter within the entire SN~Ia ejecta even before maximum light (which occurs around the 20th day after explosion), see e.g. \\cite{sorokina.EPThn} or  Fig.~2 in \\cite{sorokina.SBring02}. In this case one cannot ascribe to radiation the gas temperature, or any other temperature, since SN~Ia spectrum differs strongly from a blackbody one. One has to solve a system of time-dependent transfer equations in many energy groups instead, with an accurate prescription for treatment of a huge number of spectral lines, which are the main source of opacity in this type of SN \\cite{sorokina.BHM,sorokina.PE00}. ",
        "conclusions": "The new 3D SN~Ia model MR \\cite{sorokina.martin} is very appealing. Yet it is not a final one: a detailed post-processing of nucleosynthesis changes the composition. It has been  done very recently  \\cite{sorokina.claudia}, and it is not yet checked in the light curve calculation. Our light curve computations are also preliminary, since more work is needed on the expansion opacity. Hopefully, none of the required improvements will spoil the light curve of this model and its X-ray spectra on the SNR stage, since the specific qualities of the model can be primarily explained by the enrichment of the outermost layers of SN ejecta by  Fe and Ni.  The SN light curve modelling still has a  lot of physics to be added, such as a 3D time-dependent radiative transfer, including as much as possible of NLTE effects, which are especially essential for SNe~Ia  \\cite{sorokina.hoeflich}. All this will improve our understanding of thermonuclear supernovae.   \\subsection*{Acknowledgements}  The authors are grateful to Wolfgang Hillebrandt, to the organizers of the conference and to the MPA staff for all their help, to Stan Woosley for his hospitality at UCSC where a major part of the work on expansion opacity was done, and to Bruno Leibundgut for providing us with the data \\cite{sorokina.CLV} in electronic form.  The work is supported in Russia by RFBR grant 02-02-16500, in the US, by NASA grant NAG5-8128, and in Germany, by MPA visitor program."
    },
    "0212/astro-ph0212461_arXiv.txt": {
        "abstract": "We have developed a numerical model for the temporal evolution of particle and photon spectra resulting from nonthermal processes at the shock fronts formed in merging clusters of galaxies. Fermi acceleration is approximated by injecting power-law distributions of particles during a merger event, subject to constraints on maximum particle energies. We consider synchrotron, bremsstrahlung, Compton, and Coulomb processes for the electrons, nuclear, photomeson, and Coulomb processes for the protons, and knock-on electron production.  Broadband radio through $\\gamma$-ray light curves radiated by nonthermal protons and primary and secondary electrons are calculated both during and after the merger event. Using ROSAT observations to establish typical parameters for the matter density profile of clusters of galaxies, we find that merger shocks are weak and accelerate particles with relatively soft spectra.  Our results suggest that only a minor contribution to the diffuse extragalactic $\\gamma$-ray background can originate from cluster merger shocks. ",
        "introduction": "In the hierarchical merging cluster scenario, clusters grow in mass by accreting nearby clusters.  Approximately 30--40\\% of galaxy clusters show evidence of substructure in both the optical \\cite{beers:82} and X-ray wavelengths \\cite{forman:81}.  Velocity differences between the observed structures is $\\approx\\!\\!1000$--$2000$\\kms.  With gravitational forces driving the interaction between the two systems, cluster mergers are consistent with highly-parabolic orbits.  Typical sound speeds within the intracluster medium (ICM) are $\\approx\\!\\!1000$\\kms, so shocks will form at the interaction boundary of the two systems.  Shock fronts that form in the ICM as a result of a cluster merger event are thought to be associated with the cluster {\\em radio relics}, which are diffuse emission found on the cluster periphery with no known optical counterpart.  The shock compression will orient any existing cluster magnetic field into the plane of the shock.  Radio relics are characterized by highly organized magnetic fields with field strengths in the $\\sim\\!\\!1$\\mug\\ range with linearly polarized field lines in the vicinity of the shock \\cite{ensslin:98}.  The shock front will accelerate a fraction of the thermal particles within the ICM by first-order Fermi acceleration.  It has been proposed that cluster mergers are the dominant contributor to the diffuse $\\gamma$-ray background \\cite{loeb:00}.  Some unidentified EGRET sources are claimed to be associated with $\\gamma$-ray emission from galaxy clusters \\cite{totani:00}.  Excess EUV emission from Coma can be explained by nonthermal electrons accelerated at merger shocks \\cite{atoyan:00}. Variations in radio surface brightness will result from superposition of cluster emissions \\cite{waxman:00}. ",
        "conclusions": ""
    },
    "0212/hep-th0212326_arXiv.txt": {
        "abstract": "Classical geometry of de Sitter spacetime is reviewed in arbitrary dimensions. Topics include coordinate systems, geodesic motions, and Penrose diagrams with detailed calculations. ",
        "introduction": "Cosmological constant $\\Lambda$ was originally introduced in general relativity with the motivation to allow static homogeneous universe to Einstein equations in the presence of matter. Once expansion of the universe was discovered, its role turned out to be unnecessary and it experienced a checkered history. Recent cosmological observations suggest some evidences for the existence of a positive cosmological constant~\\cite{Rie}~: From a variety of observational sources, its magnitude is $(G\\hbar/c^{3})\\Lambda\\sim 10^{-123}$, and both matter $\\Omega_{\\rm matter}$ and vacuum $\\Omega_{\\Lambda}$ are of comparable magnitude, $\\Omega_{\\rm matter}/\\Omega_{\\Lambda}\\sim 0.3/0.7$ in spatially flat universe. Both results are paraphrased as old and new cosmological constant problems~\\cite{old,new}. Inclusion of quantum field theory predicts naturally the energy density via quantum fluctuations even up to the Planck scale, so the old problem is to understand why the observed present cosmological constant is so small. It is the worst fine-tuning problem requiring a magic cancellation with accuracy to 120 decimal places. The new problem is to understand why the vacuum energy density is comparable to the present mass density of the universe. Though there are several classes of efforts to solve the cosmological constant problem, e.g., cancellation of vacuum energy based on symmetry principle like supersymmetry, the idea of quintessence, and the anthropic principle, any of them is not widely accepted as the solution, yet. In the early universe, existing cosmological constant also provides an intriguing idea, inflationary cosmology, to fix long-standing cosmological problems in the Big-bang scenario.  In addition to the above, detailed properties of the de Sitter geometry in arbitrary dimensions raise a debate like selection of true vacuum compatible with its scale symmetry in relation with trans-Planckian cosmology~\\cite{tran,covac} or becomes a cornerstone of a theoretical idea like dS/CFT for a challenging problem of quantum gravity in de Sitter spacetime~\\cite{dscft}.  In the context of general relativity, classical study on the cosmological constant is to understand geometry of de Sitter spacetime (dS${}_d$). Therefore, in this review, we study de Sitter geometry and the motion of classical test particle in the background of the de Sitter spacetime in detail. Topics include useful coordinate systems, geodesic motions, and Penrose diagrams. One of our purposes is to combine knowledges of classical de Sitter geometry, scattered with other subjects in the textbooks or reviews, in one note~\\cite{BD,KT,text,revi}. We tried to make our review self-contained and a technical note by containing various formulas in the appendix.  The review is organized as follows. In section 2, we introduce our setup with notations and signature. In section 3, four coordinates (global, conformal, planar, and static) and coordinate transformations among those are explained with the reason why such four of them are frequently used. Killing symmetries are also obtained in each coordinate system. All possible geodesic motions of classical test particle are given in section 4. Identification of de Sitter horizon for a static observer and cosmological evolution to a comoving observer are also included. In section 5, causal structure of the global de Sitter spacetime is dealt through Penrose diagrams, and its quantum theoretical implication is also discussed. Section 6 is devoted to brief discussion. Various quantities are displayed in Appendix for convenience.  \\setcounter{equation}{0} ",
        "conclusions": "In this review, we have discussed classical geometry of $d$-dimensional de Sitter spacetime in detail~: Sections include introduction of four representative coordinates (global (or closed), conformal, planar (or inflationary), static), identification of Killing vectors, geodesic motions, cosmological implication, and Penrose diagrams.  Two important subjects are missing~: One is group theoretic approach of the de Sitter spacetime and the other is energy because we do believe they have to be systematically dealt with quantum theoretic topics like supersymmetry and Hamiltonian formulation of quantum theory. Once we try to supersymmetrize de Sitter group, inconsistency is easily encountered, that either a supergroup including both unitary representations and isometries of the de Sitter spacetime is absent or any action compatible with de Sitter local supersymmetry contains vector ghosts~\\cite{SUSY}. In de Sitter spacetime two definitions are known to define the mass of it~: One is the Abbott-Deser mass given as the eigenvalue of zeroth component of Virasoro generators and the other is the quasilocal mass obtained from Brown-York stress tensor associated to a boundary of a spacetime~\\cite{ener}. However, unitarity and conserved energy are not satisfied simultaneously in quantum theory in the de Sitter spacetime, and structure of de Sitter vacuum seems to ask more discussion~\\cite{BD,covac,KC}.  Strikingly enough recent cosmological observations suggest the existence of extremely-small positive cosmological constant, however it just leads to a new form of terrible cosmological constant problem. In the present stage, it seems to force us to find its solution by understanding quantum gravity (in de Sitter spacetime). Obviously the topics we reviewed form a basis for the following research subjects still in progress, e.g., de Sitter entropy and thermodynamics, dS/CFT correspondence, trans-Planckian cosmology, understanding of vacuum energy in the context of string theory, and etc, of which ultimate goal should be perfect understanding of the cosmological constant problem as consequence of constructed quantum gravity.    \\appendix  \\setcounter{equation}{0}"
    },
    "0212/astro-ph0212542_arXiv.txt": {
        "abstract": "In this paper we have compiled two new open cluster catalogues. In the first one, there are 119 objects with ages, distances and metallicities available, while in the second one, 144 objects have both absolute proper motion and radial velocity data, of which 45 clusters also with metallicity data available.  Taking advantages of the large number of objects included in our sample, we present an iron radial gradient of about $-$0.063$\\pm$0.008 \\dexkpc \\ from the first sample, which is quite consistent with the most recent determination of oxygen gradient by nebulae and young stars, which is about $-$0.07 \\dexkpc \\ . By dividing clusters into age groups, we show that iron gradient was steeper in the past, which is consistent with the recent result from Galactic planetary nebulae data, and also consistent with the inside-out galactic disk formation scenarios. Based on the cluster sample, we also discussed the metallicity distribution, cluster kinematics and space distribution. A disk age-metallicity relation could be implied from those properties, although we could not give conclusive result from the age metallicity diagram based on the current sample. More observations are needed for metal poor clusters. From the second catalogue, we have calculated the velocity components in cylindrical coordinates with respect to the GSR for 144 open clusters. The velocity dispersion of the older clusters are larger than that of young clusters, but they are all much smaller than that of the Galactic thick disk stars. ",
        "introduction": "Since the seminal work of \\citet{els62}, great progress has been made in understanding the formation and evolution of the Milky Way galaxy. The progress comes, on one side, from observations concerning chemical abundances in stars (and clusters), gas clouds and, on the other side, from improved knowledge relevant to galaxy formation and evolution.  However, some important quantities related to the chemical evolution of our Galaxy, such as star formation history, initial mass function, gas flow, etc., are not yet well understood. Observational data from the Milky Way disk and halo has put strong constraints on our understanding of those quantities. Among a variety of observables, radial abundance gradients along the Galactic disk is one of the most important constraints on the Galactic chemical evolution model. The existence of such gradients is now well established, through radio and optical observations of HII regions, disk stars, planetary nebulae \\citep{hen99,hou00,chi01,ma02b} and open clusters \\citep{fri95,fri99}. An average gradient of $\\sim-$0.06 \\dexkpc \\ is observed in the Milky Way disk for most of the elements, e.g. O, S, Ne, Ar and Fe. This magnitude of the observed gradients constrains the various parameters in the chemical evolution model, such as the timescales of star formation and infall \\citep{pra95} or any variations of the stellar initial mass function properties with metallicities \\citep{chi01}.  In the last decade, a number of successful models have been developed related to the chemical evolution of the Milky Way galaxy, but some important differences exist. One of them concerns the history of abundance gradient along the Galactic disk: were they steeper or flatter in the past? Different predictions were made by various models, although most of them claim that they could reproduce the majority of the observational properties both in the solar neighborhood and on the whole disk. Time flattening evolution is suggested by the models of \\citet{pra95,mol97,all98,boi99,cha99,hou00,cha02}, while the opposite is supported by models of \\citet{tos88,sam97} and \\citet{chi97,chi01}.  The situation is neither settled observationally. Estimated ages of various types (PNI, PNII, PNIII) of planetary nebulae(PN) span a large fraction of the age of the Galaxy. Observations of the abundances of those objects across the Milky Way disk could, in principle, provide some information on the time evolution of the abundance gradient \\citep{mac94,mac99,ma02a,ma02b}. In a recent work, \\citet{hou00} have made a detailed analysis for the O, Ne, S and Ar gradient based on the PN data of \\citet{mac99}. It was shown that there is fairly good agreement between model predictions and observations concerning all the properties of the observed abundance profiles (absolute values, gradient, scatter) for O, S, Ne and Ar. The model suggests that abundance gradients are steeper in the earlier epoch. However, the large scatter in the adopted data does not allow one to conclude on the temporal variation of the gradients. Nevertheless PNs suffer from large uncertainties concerning their progenitor's masses and lifetimes as well as their distances from Galactic center.  On the other hand, open clusters(OCs) have long been used to trace the structure and evolution of the Galactic disk \\citep{fri95}. Since open clusters could be relatively accurately dated and we can see them to 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 Age-Metallicity Relation(AMR), abundance gradient evolution, disk age and so on \\citep{car98}.  At this point, one might ask whether the field disk populations are also able to trace the disk evolution. Indeed, the extensive studies by \\citet{edv93}, and recently by \\citet{che00}, who concentrate on disk F, G stars, show an overall radial gradient that is nearly independent of age. Those results are based on stars mainly restricted in the solar neighborhood. A more detailed analysis for the disk iron gradient was given by \\citet{cui00} on the basis of 1302 field star with high resolution proper motion and parallax data from Hipparcos satellite. They have derived an radial iron gradient of $-$0.057 \\dexkpc \\ within galactocentric distance from 8.5 kpc to 17 kpc. However, it is still difficult to reveal any pronounced gradient evolution from those results. Moreover, results from those studies are strongly affected by selection effects and rely on the techniques for determining individual stellar distances (which are heavily dependent on the adopted Galaxy potential model) that are much less reliable than those used to obtain cluster distances. In a recent work, \\citet{cor01} have modelled the effects of the orbital diffusion of stars and clusters on the Galactic abundance gradient. The general conclusion is that the effect of diffusion makes a gradient shallower over time, and the cluster population offers a more viable means for finding detailed structure within the recent Galactic abundance gradient.  Here, we also pointed out that our recent treatment on deriving the abundance gradient from open clusters in \\citet{hou02} is in fact not proper. In that paper, we have simply taken four Catalogs from literatures \\citep{car98, twa97, pia95, fri95}, and merge them just by making cross checking for the common clusters, without examining individually to see if there are important difference among clusters in the different catalogs \\citep{twa02}. We refer this paper to act as a substitution to our previous one.  In this paper, we compiled a set of new open cluster catalogues. The catalogue was divided into two parts: CAT 1 and CAT 2. In CAT 1, we list 119 clusters with iron abundance, age, distance and reddening data available. This could provide statistically more significant information to the Galactic disk formation and evolution, such as Age-Metallicity Relation, abundance gradient and its time or/and spatial evolution and so on. The second sample consists of 144 clusters with three dimensional kinematics information available. From this sample, we are able to explore some statistical relations among kinematics and other observables.  The paper is organized as the following: firstly in Section 2, we describe the main characteristic of the two samples. Then, in Section 3, we give some statistical analysis for the sample, mainly some metallicity and kinematics distributions. The abundance gradient is given in Section 4. We show that, based on our open cluster data, the abundance gradient of the Galactic disk was steeper in the past. In Section 5, we make some detailed discussions about the Age-Metallicity Relation (AMR) of the disk. Finally, a brief summary is given in Section 6. ",
        "conclusions": "The main work of this paper is to compile a most complete open clusters sample with metallicity, age, distance data as well as kinematic information available. And upon this sample, some statistical analysis on spatial and metallicity distributions have been made.  We derived an iron radial gradient about $-$0.063$\\pm$0.008 \\dexkpc \\ from the CAT 1, which is quite consistent with the most recent determination of oxygen gradient in nebulae and young stars. By dividing clusters into age groups, we show that iron gradient was steeper in the past, which is consistent with the recent result from Galactic planetary nebulae data. Our result supports the inside-out Galactic disk formation mechanism that invoking SFR and infall time scale various with radius.  The spatial variation of gradient was also explored. When the clusters were divided into inner and outer groups, we found that the abundance gradient was a bit shallower inside 10 kpc. But the inner most cluster in our sample is in $R_{GC} = 6.8$kpc, we need more data of clusters in the inner region. This could be helpful to judge the radial flow effect on the current galactic chemical evolution model.  From scale heights of metal poor and metal rich clusters, we noticed that the metallicity could be related to the age. However, by plotting directly the dependence of cluster abundancess on their ages, no strike slope in AMR was found. However, the paucity of metallicity of very old open clusters made it impossible to give a definite conclusion based on the current sample."
    },
    "0212/astro-ph0212297_arXiv.txt": {
        "abstract": "We analyze the topology, lifetime, and emissions of a torus around a black hole formed in hypernovae and black hole-neutron star coalescence. The torus is ab initio uniformly magnetized, represented by two counter oriented current-rings, and develops a state of suspended accretion against a ``magnetic wall\" around the black hole. Magnetic stability of the torus gives rise to a new fundamental limit ${\\cal E_B}/{{\\cal E}_k}<0.1$ for the ratio of poloidal magnetic field energy-to-kinetic energy, corresponding to a maximum magnetic field strength $B_c\\simeq 10^{16}\\mbox{G}\\left({7M_\\odot}/{M_H}\\right) \\left({6M_H}/{R}\\right)^2\\left({M_T}/{0.03M_H}\\right)^{1/2}$. The lifetime of rapid spin of the black hole is effectively defined by the timescale of dissipation of black hole-spin energy $E_{rot}$ in the horizon, and satisfies $T\\simeq 40\\mbox{s} (M_H/7M_\\odot)(R/6M_H)^4(0.03M_H/M_T)$ for a black hole of mass $M_H$ surrounded by a torus of mass $M_T$ and radius $R$. $E_{rot}$ of the black hole. The torus converts a major fraction $E_{gw}/E_{rot}\\sim 10\\%$ into gravitational radiation through a finite number of multipole mass-moments, and a smaller fraction into MeV neutrinos and baryon-rich winds. At a source distance of 100Mpc, these emissions over $N=2\\times 10^4$ periods give rise to a characteristic strain amplitude $\\sqrt{N}h_{char}\\simeq 6\\times 10^{-21}$. We argue that torus winds create an open magnetic flux-tube on the black hole, which carries a minor fraction $E_j/E_{rot}\\simeq 10^{-3}$ in baryon-poor outflows to infinity. We conjecture that these are not high-sigma outflows owing, {in part}, to magnetic reconnection in surrounding current sheets. The fraction $E_j/E_{rot}\\sim (1/4)(M_H/R)^4$ is standard for a universal horizon half-opening angle $\\theta_H\\simeq M_H/R$ of the open flux-tube. We identify this baryon poor output of tens of seconds with GRBs with contemporaneous and strongly correlated emissions in gravitational radiation, conceivably at multiple frequencies. Ultimately, this leaves a black hole binary surrounded by a supernova remnant. ",
        "introduction": "Black holes surrounded by a magnetized torus or disk are believed to constitute the central engines that power various high-energy sources, notably active galactic nuclei, galactic microquasars, and gamma-ray bursts. The latter systems are thought to be the outcome of catastrophic events such as core-collapse in massive stars and black hole-neutron star coalescence \\citep{eich89,woo93,pac91,pac98}, and are of interest as potentially the most extreme and short-lived black hole-torus systems.  We describe a theory for the topology, lifetime and emissions of a black hole-torus sytems as a function of three parameters: the mass $M_H$ of an extreme Kerr black hole, the radius $R$ and the mass $M_T$ of the torus. We shall do so largely by {studying the torus} by equivalence to pulsars, {in both topology and ms rotation periods.} The energy emissions are powered by the spin-energy of a Kerr black hole \\citep{ker63}. Most of the black hole-luminosity {-- the rate at which the black hole deposits energy into its surroundings in all channels --} is incident on the torus, which hereby creates a {\\em major} energy output of the system. A {\\em minor} energy output is released in baryon-poor outflows through an open magnetic flux-tube along the spin-axis of the black hole. The life-time of these black hole-torus systems is identified with the lifetime of rapid spin of the black hole in a state of suspended accretion \\citep{mvp01A} against a ``magnetic wall\" around the black hole \\citep{mvp99}. The suspended accretion state results from a strong coupling of the torus to the spin-energy of the black hole. We point out that this mechanism is based on a uniform magnetization of the torus, represented by two oppositely oriented current rings \\citep{mvp99}. This magnetization is a natural outcome of both black hole-neutron star coalescence and core-collapse in hypernovae.  In this paper, we quantify (1) {the lifetime of rapid spin of the black hole in terms of a new magnetic stability criterion for the torus}; (2) baryon-rich outflows from the torus at MeV temperatures; and (3) the fraction of black hole spin-energy in baryon-poor outflows through an open magnetic flux-tube on the black hole, created from outer layers of the inner torus magnetosphere by these torus winds.  The torus is luminous in various channels, which are strongly correlated by the properties of the torus: in gravitational radiation, winds, thermal and MeV neutrino emissions \\citep{mvp01a,mvp02}. Calorimetric constraints on the torus winds hereby obtains predictions for the proposed emissions in gravitational radiation, while calorimetry on the gravitational wave emissions by upcoming gravitational wave-experiments obtains a method for identifying Kerr black holes as objects in nature \\citep{mvp02}.  Gravitational radiation forms a major output of the system and the dominant output of the torus \\citep{mvp01a,mvp02}. This is emitted by a finite number of multipole mass moments in a torus of finite slenderness, due to the Papaloizou-Pringle instability \\citep{pap84,mvp02c} and, conceivably, other wave-modes in the torus. These emissions are candidate sources for the upcoming gravitational wave-experiments by laser {interferometric instruments LIGO \\citep{abr92}, VIRGO \\citep{bra92}, TAMA \\citep{mas01} and GEO (e.g., \\cite{geo}),} or by any of the bar or sphere detectors presently under construction. The frequency in gravitational radiation is determined by the keplerian frequency of the torus. The latter is strongly correlated to the output energy in torus winds. This suggests the possibility of performing calorimetry on the impact of these torus winds on the remnant stellar envelope \\citep{mvp02d} and on hypernova remnants, in order to constrain the expected frequency in gravitational radiation.  Baryon-poor outflows form a small fraction of the total output from the black hole through an open magnetic flux-tube \\citep{mvp02}. We here attribute the formation of this open flux-tube to powerful baryon-rich torus winds, driven from the surface by escaping MeV neutrinos. This neutrino output provides the dominant cooling mechanism of the torus during the suspended accretion state, in addition to the energy release in gravitational waves. The fraction of rotational energy thus released is proportional to $\\theta_H^4$, which is standard for a universal horizon half-opening angle $\\theta_H$ of the inner tube. These outflows are probably not high-sigma, owing to magnetic reconnection in the interface between the baryon poor and baryon-rich winds in a surrounding outer flux-tube. The half-opening angle is possibly related to curvature in poloidal topology \\citep{mvp02d}.  Tentative observational evidence for a black hole-luminosity incident into surrounding matter in case of supermassive black holes is found in MCG-6-30-15 \\citep{wil01}, and in case of stellar mass black hole-candidates in the galactic source XTE J1650-500 \\citep{mil02}. Independent observational evidence of the presence of magnetic fields remains elusive. We point out, however, that in our model heating of the torus is due to viscous shear between its inner and outer faces, possibly in the form of magnetohydrodynamical turbulence, which is different from electric dissipation as envisioned in \\cite{wil01}.  In \\S\\ref{sec:T-magnetosphere} and \\S\\ref{sec:stability} we describe the topology and stability of the magnetosphere of a torus formed in black hole-neutron star coalescence and hypernova, the suspended accretion state around rapidly rotating black holes and the secular time-scale {of its rapid spin}. The formation of a finite number of multipole mass moments by the Papaloizou-Pringle instability in tori of finite slenderness is summarized in \\S\\ref{sec:susaccrt}. Energy emissions in various channels by the torus are described in \\S\\ref{sec:dissp}. We calculate the neutrino driven mass loss rate from the torus' surface in \\S\\ref{sec:massejection}. These baryon-rich torus winds have some implications for energy extraction from the black hole and the creation of open magnetic flux-tubes from outer layers in the inner torus magnetosphere. \\S \\ref{sec:opentubes}- \\S \\ref{Sec:interface} describe the structure of the inner flux-tube supported by the black hole, and dissipation by magnetic reconnection in its interface with the surrounding outer flux-tube. We conclude with observational consequences in \\S\\ref{Sec:conclusion}. ",
        "conclusions": ""
    },
    "0212/astro-ph0212404_arXiv.txt": {
        "abstract": "{ Radial and vertical profiles are determined for a sample of 34 edge-on disk galaxies in the HDFs, selected for their apparent diameter larger than 1.3\\arcsec and their unperturbed morphology. The thickness and flatness of their galactic disks are determined and discussed with regard to evolution with redshift. We find that sub-$L^*$ spiral galaxies with $z \\sim 1$ have a relative thickness or flatness (characterized by $h_z/h$ the scaleheight to scalelength ratio) globally similar to those in the local Universe. A slight trend is however apparent, with the $h_z/h$ flatness ratio larger by a factor of $\\sim1.5$ in distant galaxies if compared to local samples. In absolute value, the disks are smaller than in present-day galaxies. About half of the $z \\sim 1$ spiral disks show a non-exponential surface brightness distribution. ",
        "introduction": "A very fundamental characteristic of galactic disks is the flatness, defined by the ratio between exponential scalelengths in the vertical and radial direction  $h_z/h$. The determination of the thickness of stellar disks is a difficult enterprise and has been determined mostly for edge-on galaxies, although attempts have been made in other cases, too (Ma et al. 1998). Typical values are found to be between 0.15 and 1 (van der Kruit \\& Searle 1981, 1982; de Grijs 1998) with a systematic trend for galaxies to become thinner from S0 to Sc if a more complex multicomponent structure (thick/thin disk) is not taken into account. Stellar velocity dispersion measurements show that the dispersion tends to grow with total luminosity, and vary radially as the square root of surface density, so as to support a constant scaleheight (Bottema 1993). Recently Kregel et al. (2002) have shown that the flattening of disks increases with the amplitude of their maximum rotation, and their total HI mass.  This flatness is an important clue to the formation and evolution history of galaxies. Disk stars are thought to be formed out of a very thin gaseous disk. The observed structure then thickens due to scattering by local fluctuations of the gravitational field as caused, e.g., by giant molecular clouds or  spiral structure in the disk. The velocity dispersion of the stars increases in both the radial and vertical direction by this diffusion through phase space (Wielen 1977, Binney \\& Lacey 1988). The disk may thicken more efficiently  through interactions with small companions or minor mergers, as suggested by numerical simulations (Quinn et al. 1993, Walker et al, 1996, Velazquez \\& White 1999). The thickening could amount to a factor of 1.5--2 and can be a constraint on the frequency of mergers if disks are observed today to be too thin (e.g. Toth \\& Ostriker 1992). The existence of a thick stellar disk in our own Galaxy has been attributed to an old merger event (Robin et al 1996, Dalcanton \\& Bernstein 2002). However, other phenomena have to be taken into account since the thin disk could also be maintained through gas infall with subsequent star formation. Many observational facts strongly suggest a high gas infall rate (e.g. Toomre, 1990, Sancisi et al. 1990, Jiang \\& Binney 1999). Dynamics of galaxy disks, bar and spiral reformation constrain this rate such that galaxies may double their mass in less than 10 Gyrs (Bournaud \\& Combes 2002, Block et al 2002).  The influence of tidal interaction on the flatness of galaxy disks  has been addressed by Reshetnikov et al. (1993) and Reshetnikov \\& Combes (1997) who found in interacting galaxies  the ratio $h_z/h$ to be a factor of two  higher, due to a disk thicker in absolute values, but also shorter in radial dimensions. Schwarzkopf \\& Dettmar (2000, 2001) confirmed this trend on a larger sample, but essentially because of a larger absolute thickness of interacting disks. They also noted that tidally interacting galaxies are more frequently warped and that the thickening was more noticeable in the outer parts.  All these studies suggest that the relative flatness of a galaxy disk can vary by both effects during its evolution since the majority of galaxies have experienced interactions in the past. It appears that galaxies were dynamically ``hotter'' in the past (Abraham et al. 1999), which is seen in a smaller fraction of barred galaxies, and this could also have some consequences for the flattening. As the frequency of interactions/mergers strongly  increases with redshift (e.g. Le F\\`evre et al 2000), it is interesting to trace the evolution of the disk flatness with time by studying photometrically distant galaxies with high spatial resolution. This is now possible in the Hubble Deep Fields HDF-N and HDF-S and we present below the results of such a study. Section 2 describes our sample and Sect. 3 gives the  derived general characteristics of disks in terms of radial and vertical profiles. The last section discusses the main results on the evolution of the flatness for  stellar disks.  Throughout the paper, we adopt a flat cosmology with $\\Omega_0=1$ and $H_0 = 70$ km~s$^{-1}$ Mpc$^{-1}$. (For a model with $\\Omega_m=1/3$, $\\Omega_{\\Lambda}=2/3$ and $\\Omega_m+\\Omega_{\\Lambda}=1$, linear sizes at $z=0.7-1$ will be $\\approx$25\\% larger and absolute magnitudes are brigther by $\\approx$0\\fm45.) If not differently specified, all magnitudes are expressed in the AB system (Oke 1974). ",
        "conclusions": "Our sample of distant edge-on galaxies shows an unexpectedly large fraction of non-exponential disks. From this observation we infer that many sub-$L^*$ disks are still in the formation process at $z \\sim 1$, or the degree of tidal perturbations is higher, which has already been observed (e.g. Le F\\'{e}vre et al. 2000). One of the possible mechanisms to redistribute matter and to develop the exponential profile is the action of a long-lived bar (e.g. Combes et al. 1990). The young galaxies of our sample may not have experienced enough bar torques for this process, either because of very transient bars in very unstable disks, or because the effects of bars are continuously perturbed by mergers and accretion. The consequence is that a galaxy will spend much less of its time in a barred state. This conclusion is in good agreement with the general depletion of barred spirals at $z > 0.7$ (Abraham et al. 1999, van den Bergh 2002).  We have taken into account the effect of the PSF to demonstrate that the distant disks examined in this work are thicker by a factor of 1.5 if compared to  local disk sample. This is a preliminary result based on a small number of sub-$L^*$ galaxies, and should be confirmed by larger statistical samples.  Since the star formation rate increases strongly with redshift (e.g. Madau et al. 1996, Flores et al. 1999, Thomson et al. 2001) one could  speculate that the fraction of young stars formed in the midplane of disks is expected to be much higher in the past, making the young disk appear much thinner. Previous work on local samples has already concluded that tidal interactions between galaxies have a certain thickening effect of the order of 1.5-2 (Reshetnikov \\& Combes 1997, Schwarzkopf \\& Dettmar, 2000). Since the frequency of galaxy interactions/mergers increases strongly with redshift, the present result fits into this picture. If one assumes that all disks have undergone such interactions, either additional cooling processes for the hot disks or a substantial growth predominantly in the radial direction are required to explain the flatness of disks in the local Universe.  However, other parameters should be taken into account in the interpretation. The distant disks are on absolute scales smaller than the local sample disks. It is likely that the morphological types are later (less evolved) and the masses of the galaxies smaller on average, as already found in previous studies based on HST surveys (e.g. Abraham et al. 1996, Glazebrook et al. 1998). The lower frequency of bars, associated with the smaller bulge-to-disk ratio, suggests that disks are more unstable and dynamically hotter (Abraham et al. 1999), which could explain the present thickening effect. In this context it is also noteworthy to mention that the galaxies of our sample contain only a few bulges as can be judged from the major axis profiles in Fig.~8. This can either be interpreted as evidence for a later formation of bulges by secular processes (e.g. Combes 2001) or that our sample is selected such that it represents early phases of late morphological types only."
    },
    "0212/astro-ph0212318_arXiv.txt": {
        "abstract": "We present deep Near--Infrared (NIR) imaging of Blue Compact Dwarf Galaxies (BCDs), allowing for the first time to derive and systematize the NIR structural properties of their stellar low--surface brightness (LSB) host galaxies. Compared to optical data, NIR images, being less contamined by the extended stellar and ionized gas emission from the starburst, permit to study the LSB host galaxy closer to its center. We find that radial surface brightness profiles (SBPs) of the LSB hosts show at large radii a mostly exponential intensity distribution, in agreement with previous optical studies. At small to intermediate radii, however, the NIR data reveal an inwards flattening with respect to the outer exponential slope (``type V SBPs'', Binggeli \\& Cameron 1991) in the LSB component of more than one half of the sample BCDs. This result may constitute an important observational constraint to the dynamics and evolution of BCDs.  We apply a modified exponential fitting function (Papaderos et al. 1996a) to parametrize and systematically study type V profiles in BCDs. A S\\'ersic law is found to be less suitable for studying the LSB component of BCDs, since it yields very uncertain solutions. ",
        "introduction": "Studies of Blue Compact Dwarf Galaxies (BCDs) are crucial for understanding the starburst--driven evolution of gas--rich low--mass extragalactic systems at low and high redshift. Particularly important for the study of BCDs is the information on their stellar low-surface brightness (LSB) host galaxy (Loose \\& Thuan \\citeyear{loose86}). This evolved component, present in most types of local dwarf galaxies, contains the bulk of the stellar mass of a BCD and is most probably dynamically relevant. It is therefore likely that its structure (e.g., mass density distribution) influences the global star formation process in a BCD. However, the intensity distribution of the LSB component is uncertain close to its center, where it is most required to understand the centrally concentrated star--forming (SF) activity of typical BCDs. Previous optical studies of these systems have been restricted to the LSB periphery ($\\gsim$2 exponential scale lengths) where the emission of the starburst becomes negligible. At NIR wavelengths, the fractional flux contribution of the starburst is comparatively low, allowing to study the LSB hosts of BCDs closer to their center. We present results from the first large BCD sample, comprising $>$ 30 objects of all mor\\-pho\\-lo\\-gi\\-cal sub\\-clas\\-ses, which was observed sufficiently deep in the NIR to sys\\-te\\-ma\\-ti\\-cal\\-ly study the structural properties of the LSB hosts. The complete analysis of the sample is presented in Noeske et al. (\\citeyear{noeske02}) and Cair\\'os et al. (\\citeyear{cairos02}). ",
        "conclusions": ""
    },
    "0212/astro-ph0212154_arXiv.txt": {
        "abstract": "We present new arcminute resolution radio images of the low surface brightness radio source PKS~B1400$-$33 that is located in the poor cluster Abell S753. The observations consist of 330 MHz VLA, 843 MHz MOST and 1398 and 2378 MHz ATCA data. These new images, with higher surface brightness sensitivity than previous observations, reveal that the large scale structure consists of extended filamentary emission bounded by edge-brightened rims. The source is offset on one side of symmetrically distributed X-ray emission that is centered on the dominant cluster galaxy NGC 5419. PKS~B1400$-$33 is a rare example of a relic in a poor cluster with radio properties unlike those of most relics and halos observed in cluster environments.  The diffuse source appears to have had an unusual origin and we discuss possible mechanisms.  We examine whether the source could be re-energized relic radio plasma or a buoyant synchrotron bubble that is a relic of activity in NGC 5419. The more exciting prospect is that the source is relic plasma preserved in the cluster gaseous environment following the chance injection of a radio lobe into the ICM as a result of activity in a galaxy at the periphery of the cluster. ",
        "introduction": "Extended emission structures with steep radio spectral indices, which are not obviously associated with any optical galaxy, are sometimes observed in rich cluster environments.  These sources are broadly separated into cluster wide `halos' that appear relaxed and symmetric with respect to the X-ray intracluster gas, and the peripheral arc-like structures referred to as `relics'.   Both types are predominantly observed to be associated with hot ($T_{X} \\ge 6$ keV), X-ray luminous ($L_{X} > 4\\times 10^{44}$ erg s$^{-1}$) rich compact clusters \\citep{schuecker99}; however, the clusters that have peripheral relics may have somewhat lower temperatures \\citep{feretti99}.  A standard hypothesis is that these sources are relic synchrotron plasma that has been revived.  PKS~B1400$-$33 is an extended radio source, with an 85 MHz flux density of 57 Jy \\citep{mills60}, and has the lowest surface brightness of any source in the Parkes catalog. PKS~B1400$-$33 appears to be associated with the poor cluster S753 of Abell richness class 0; however, the source not been conclusively identified with any optical galaxy.  Previous low frequency images of the source using the VLA and the MOST \\citep{goss87} showed a relatively bright rim along the NE edge and low surface brightness filamentary emission trailing off to the SW.    The source was observed to have a steep spectrum with spectral index $\\alpha$ ($S_{\\nu} \\propto \\nu^{-\\alpha}$) in the range 1.2-2.4.  The unusual nature of this possible relic radio source and its potential value as a probe of the gas dynamics and evolution of poor cluster environments has led us to carry out new radio observations of the source.  We have made images with improved surface brightness sensitivity at 330 MHz using the Very Large Array (VLA), at 843 MHz with the Molonglo Observatory Synthesis Telescope (MOST) and at 1398 and 2378 MHz using the Australia Telescope Compact Array (ATCA).  Table 1 is a summary of the new radio observations.  These are presented here followed by discussions on the origin of this source in the light of the current understanding of the phenomenology and formation of relics in the intracluster medium (ICM). ",
        "conclusions": "We have presented new and improved radio images of the relic source PKS~B1400$-$33.  We have discussed the origin of this relic in the light of recent progress in the understanding of relic synchrotron plasma in cluster environments.  The diffuse source has an extremely low surface brightness and a steep spectrum, with radio properties unlike that of relics and halos observed in cluster environments.  The unusual morphology, placement in the cluster and physical parameters indicate an unusual origin for this source.  The new data does not represent evidence suggesting that the relic source PKS~B1400$-$33 was created by past activity in NGC 5419.  Nor do the observations rule out this possibility.  The discussions suggest that PKS~B1400$-$33 is not a typical relic.  The radio properties suggest that if the source was created by processes similar to those that form the relics and halos, we might be observing an extreme form of a relic in a poor cluster environment or, perhaps, a relic of a re-energized relic.  We have presented marginal evidence supportive of the hypothesis that the diffuse source PKS~B1400$-$33 was born as a lobe of a powerful radio galaxy; however, we do not have a good candidate for the optical host.  The source is unusual and the peculiar radio properties might be indicating an unusual origin: we favour the interpretation that the diffuse source is a relic of a lobe injected into the cluster environment as worthy of further consideration, particularly because of the novelty of the proposed phenomenon. Followup low frequency imaging of the field, possibly with the VLA at 74 MHz, with the aim of confirming the MOST detection of a counterlobe and better quality imaging of any low surface brightness connecting features, is proposed as the next step towards understanding the phenomenology in this unusual source."
    },
    "0212/astro-ph0212362_arXiv.txt": {
        "abstract": "% We have written an automatic image processing pipeline for the Advanced Camera for Surveys (ACS) Guaranteed Time Observation (GTO) program.  The pipeline, known as Apsis, supports the different cameras available on the ACS instrument and is written in Python with a flexible object-oriented design that simplifies the incorporation of new pipeline modules. The processing steps include empirical determination of image offsets and rotation, cosmic ray rejection, image combination using the drizzle routine called via the STScI Pyraf package, object detection and photometry using SExtractor, and photometric redshift estimation in the event of multiple bandpasses. The products are encapsulated in XML markup for automated ingestion into the ACS Team archive. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/astro-ph0212538_arXiv.txt": {
        "abstract": "The kinematics of galaxies with 10 megaparsecs (10 Mpc) of the Milky Way is investigated using published distances and radial velocities.  With respect to the average Hubble flow (isotropic or simple anisotropic), there is {\\em no} systematic relation between peculiar velocity dispersion and absolute magnitude over a range of 10 magnitudes; neither is there any apparent variation with galaxy type or between field and cluster members.  There are several possible explanations for the lack of variation, though all have difficulties: either there is no relationship between light and mass on these scales, or the peculiar velocities are not produced by gravitational interaction, or the background dynamical picture is wrong in some systematic way.  The extremely cold local flow of 40-60 km s$^{-1}$ dispersion reported by some authors is shown to be an artifact of sparse data, a velocity dispersion of over 100 km s$^{-1}$ being closer to the actual value.  Galaxies with a high (positive) radial velocity have been selected against in studies of this volume, biasing numerical results. ",
        "introduction": "Under the physically well-motivated assumption that the only important force on the largest astronomical scales and for most of the history of the universe is gravity, there should be an intimate relationship between the cosmic velocity field and the cosmic density field.  Predicting, calculating and observing the details of this relationship have been the major occupations of cosmology for most of the past century.  The first-order description of the kinematics of the universe takes the form of a homogeneous, isotropic expansion.  To investigate this it is best to examine the largest spatial scales, where the density field is closest to the homogeneous ideal. In that case the observational goal is to determine the Hubble constant ($H_0$) and the major problem is the determination of reliable distances.  A subsidary problem appears in practice, seen for example in \\citet{VB79}, in the fact that the distribution of galaxies with the most reliable distances (the closer ones) is not homogeneous and their motions are subsequently not uniform.  With improving observational and calculational techniques this subsidary problem has been turned into an area of research in its own right. \\citet{CW00} provide an overview; specific techniques and results are found in, for instance, \\citet{BTF99} and \\citet{DEK99}.  One of the objects of this work is the estimations of the biasing parameter, which measures the relationship between actual density fluctuations and observed galaxies.  To do this convincingly, however, requires a volume large enough that density contrasts remain in or near the linear region, hence this is the study of large-scale structure.  Unfortunately, distance indicators for galaxies so far away are not very accurate and are subject to various systematic biases; one must take large samples and be careful about manipulating statistics.  The situation reverses itself in the Local Volume, the region within about 10 megaparsecs (Mpc).  Here we are several rungs down on the cosmic distance ladder and distance indicators are based on resolved stars, allowing accuracies better than 7\\% under good conditions.  With redshifts good to a few km s$^{-1}$ from HI observations, the data are (potentially) excellent. Balancing this are the complications of nonlinear dynamics (sufficiently nonlinear that pertubative techniques, such as the Zeldovich approximation and expansion to a few orders, cannot be used) and the fact that peculiar velocities are comparable to the Hubble flow.  Dynamically, this is almost unknown territory.  It lies between the well-studied galaxy and galaxy group scales on one hand, and that of large-scale structure on the other.  Neither are there any galaxy clusters in the Local Volume, Virgo being just outside.  The potential scientific returns of examining this unexplored region, especially when excellent data are available, are great.  For instance, dark matter on galaxy scales seems to be concentric with luminous matter, but far more extended; on the linear-structure scale, both seem to be distributed the same way.  The biasing parameter $b$, which measures the concentration of luminous matter relative to dark matter, is found to be in the range 0.5-1.0 on large scales (\\citet{DEK99}; \\citet{BTF99}); within a spiral galaxy it must range from zero beyond the visible disk to some number possibly greater than unity at the center.  In between these two scales, what happens?  Without a dynamical model of obvious applicability to the Local Volume, we must begin by studying the kinematics.  The procedure in what follows is to calculate overall average motions, and then examine deviations from them for systematic effects.  In effect, we are looking for some kind of non-randomness among peculiar (radial) velocities.  Two models will be used for the ``overall average motions.''  Both include the average velocity of all galaxies in a sample, showing up as the solar reflex velocity; the first model then adds a uniform expansion (the local reflection of the universal Hubble flow).  The second is suggested by the general distribution of galaxies within the Volume. The Supergalactic Plane is well-defined in this region, clear in \\citet{TF87} as well as with updated data in \\citet{KMa96} and \\citet{LV00}.  There should be therefore an anisotropy in the overall flow, expansion normal to the Supergalactic Plane being slower than in the Plane.  In fact \\citet{VB79} found something of the sort, as have \\citet{KM01} and \\citet{KMa96}. In addition, the tidal effect of the Virgo Cluster should be visible as an elongation (high eigenvalue) in the direction of the Cluster and a contraction (low eigenvalues) normal to it.  Since Virgo is almost in the Plane, at roughly Supergalactic Longitude $100\\arcdeg$, we expect the highest eigenvalue in that direction, the lowest at the Supergalactic poles, and an intermediate eigenvalue in the Plane at longitude about $10\\arcdeg$. This relationship between the distribution of galaxies in the Supergalactic Plane and velocities can be seen either as dynamic (the concentration of mass causing anisotropic motions) or as kinematic (the concentration of galaxies having been produced by anisotropic motions), or more accurately both. ",
        "conclusions": "An examination of the kinematics of galaxies within 10 Mpc of the Milky Way has thrown up some surprises and one deep puzzle.  An overall anisotropic expansion, expected from the distribution of galaxies in the Local Volume, can be calculated; but the uncertainty of its details and its strong dependence on the particular data set chosen indicate that it is not a useful description.  A mismatch between (on one hand) the calculated standard error of various model parameters and (on the other) the change of those parameters between data sets indicates that the present data do not form a good sample of the underlying kinematics, at least in the context of the models at hand.  At least part of this is due to a selection bias against nearby galaxies with high radial velocities.  There is {\\em no} variation in the width of the peculiar velocity dispersion with absolute magnitude over a range of ten magnitudes.  Neither is there any apparent relationship with galaxy type, or between field and cluster galaxies.  This is difficult to understand, as the tendencies which even out peculiar velocities among masses on larger scales are not at work.  Mass may have no relation to light on these scales; or peculiar velocities might be produced by other than gravitational interactions among galaxies; or the set of data at hand might somehow be terribly misleading; or the reference kinematic model might be wrong.  To clarify the kinematics of this volume, and to shed light on (much less clear up) the puzzle, much additional data will be necessary.  In particular, to define the shape of the peculiar velocity histogram many more galaxies need to be added.  Unfortunately, accurate distances to objects between 2 and 10 Mpc require careful observations with the larger telescopes, and distances to a useful fraction of hundreds of galaxies will take much time to compile.  It will be worthwhile to target, initially, bright and massive galaxies;  those at high supergalactic latitude; those in the field; and those of high radial velocity but possibly still nearby."
    },
    "0212/astro-ph0212224_arXiv.txt": {
        "abstract": "We present results from an investigation of the dynamical behavior of buoyant magnetic flux rings in the radiative interior of a uniformly rotating early-type star.  Our physical model describes a thin, axisymmetric, toroidal flux tube that is released from the outer boundary of the convective core, and is acted upon by buoyant, centrifugal, Coriolis, magnetic tension, and aerodynamic drag forces.  We find that rings emitted in the equatorial plane can attain a stationary equilibrium state that is stable with respect to small displacements in radius, but is unstable when perturbed in the meridional direction. Rings emitted at other latitudes travel toward the surface along trajectories that largely parallel the rotation axis of the star.  Over much of the ascent, the instantaneous rise speed is determined by the rate of heating by the absorption of radiation that diffuses into the tube from the external medium.  Since the time scale for this heating varies like the square of the tube cross-sectional radius, for the same field strength, thin rings rise more rapidly than do thick rings. For a reasonable range of assumed ring sizes and field strengths, our results suggest that buoyancy is a viable mechanism for bringing magnetic flux from the core to the surface, being capable of accomplishing this transport in a time that is generally much less than the stellar main sequence lifetime. ",
        "introduction": "Although magnetic fields and related activity are believed to be nearly ubiquitous among stars of solar and lower  mass, it is less certain that magnetism is a common characteristic of stars more massive than the Sun.  To date, definitive detections of fields in stars with masses $\\gta 1.5\\ M_\\odot$ have, for the most part, been made for objects having anomalous chemical abundances (e.g., the chemically peculiar A and B stars; Landstreet 1992; Donati 1998).  Recently, however, observations of cyclic variability in the properties of winds from luminous OB stars have been interpreted as evidence for the presence of large-scale magnetic fields in the surface layers and atmospheres of these objects (see, e.g., Kaper \\& Henrichs 1994; Kaper et al. 1997; Prinja 1998). These inferences have been bolstered by the unambiguous measurement of a weak ($\\sim 360$ G) field in the chemically normal B1 IIIe star $\\beta$ Cephei (Donati et al. 2001).  These results suggest that magnetic fields of moderate strength might be more prevalent among hot stars than had previously been thought (see Charbonneau \\& MacGregor 2001, hereafter Paper I, for a more complete discussion).  At the present time, the origin of magnetism in massive stars is not well understood.  It is unclear whether the inferred fields are fossil in nature, or are instead the product of dynamo activity in the stellar interior (see, e.g., Parker 1979).  The latter explanation, if appropriate, raises additional interesting questions concerning the site of magnetic field generation inside hot stars. For stars having spectral types O and B, convection occurs in the innermost portion of the core, not in the outer envelope as in low-mass stars.  Because convection is thought to be necessary in order that field regeneration via the so-called $\\alpha$-effect take place, it follows that a hot star dynamo should be located deep within the stellar interior.  Within the context of mean-field electrodynamics, the properties of kinematic core dynamo models for early-type stars have been investigated by Levy \\& Rose (1974) and Sch\\\"ussler \\& P\\\"ahler (1978), and, more recently, in Paper I.  If the magnetic field of a hot star is produced by dynamo action in the convective core, then a mechanism for transporting the field to the stellar surface must be identified.  The finite electrical conductivity of the envelope leads to the outward diffusion of any fields contained therein, but only over an extended period of time.  Estimates indicate that for stars more massive than a few solar masses, the resistive diffusion time across the radiative interior exceeds the main sequence lifetime $t_{MS}$ (Sch\\\"ussler \\& P\\\"ahler 1978).  Another possibility is that dynamo fields are advected from the core to the surface by rotation-induced meridional circulation (see Paper I). For a star of mass $M_*$ and equatorial rotation speed $v_{rot}$, a rough estimate for the circulation time $t_c$ as a fraction of $t_{MS}$ is $(t_c/t_{MS}) \\approx 9 \\times 10^2\\ (M_*/M_\\odot)^{0.8}\\ (v_{rot} / 1\\ {\\rm km\\ s}^{-1})^{-2}$ (Kippenhahn \\& Weigert 1994).  In the case of a star with $M_* = 10\\ M_\\odot$ and $v_{rot} = 150$ km s$^{-1}$, $(t_c / t_{MS}) \\approx 0.25$, indicating that even for relatively rapidly rotating stars, a time $\\sim t_{MS}$ is required to bring the field produced by the dynamo to the surface by this mechanism.  In addition, the results of Paper I suggest that the circulatory flow generated by extremely rapid rotation is capable of interfering with the operation of a core dynamo.  Alternatively, in the Sun, magnetic flux emerges at the photosphere in the form of fibrils or flux tubes.  The field is thought to assume this form in or near the dynamo domain at the bottom of the convective zone.  The subsequent rise of a flux tube to the solar surface is driven by buoyancy, a consequence of the reduced density inside a tube that is in mechanical and thermal equilibrium with the surrounding, adiabatically stratified, field-free gas.  If the formation of flux tubes is a process that is not specific to fields in the Sun, then a dynamo-generated field at the base of the extended radiative envelope of a hot star might also develop fibril structure.  In this case, the buoyant force might likewise enable flux from deep in the interior to reach the surface in a time that is shorter than evolutionary time scales.  An important distinction, however, is that unlike the the solar convection zone, the stratification within the radiative interior of a hot star is sub-adiabatic.  The motion of buoyant magnetic elements in a thermodynamic environment like that in the envelope of an upper main sequence star has been studied by Gurm \\& Wentzel (1967) and Moss (1989).  In the present paper, we adopt the premise that hydromagnetic dynamo activity takes place inside the convective core of a rotating massive star, in a manner similar to that described in Paper I.  We furthermore presume that the fields so-produced naturally form into discrete, toroidal flux tubes that remain in total pressure equilibrium with the external, unmagnetized stellar interior at all times.  Using this conceptual framework in concert with the physical model described in \\S2, we determine the time-dependent position, velocity, and thermodynamic properties of a buoyant flux ring that begins its outward motion from a specified location on the core boundary.  Results pertaining to the dynamics and rise times of rings with a variety of initial field strengths and cross-sectional radii are presented in \\S3.  Our findings regarding the dynamical behavior of magnetic rings in a radiative environment and the efficacy of buoyancy as a flux transport mechanism in hot stars are summarized in \\S4. ",
        "conclusions": "The foregoing examination of the physics of magnetic flux rings in the radiative envelope of a uniformly rotating hot star has revealed a range of possible dynamical behaviors, and has enabled us to estimate the efficiency of buoyancy as a means of transporting dynamo-generated fields from the core to the surface. Rings located in the equatorial plane can attain a stationary equilibrium in which $u_r = 0,\\ u_\\phi \\neq 0$, and the combined Coriolis and tension forces balance buoyancy.  This state is stable against infinitesimal displacements of the ring in the radial direction, but unstable when the perturbations are meridionally directed, causing the ring position to shift from latitude $0^\\circ$. Apart from transient episodes of oscillatory behavior at the outset of their motion, rings released from latitudes other than $0^\\circ$ move to larger radii along paths that are very nearly parallel to the stellar rotation axis.  During most of the ascent, the rate of rise of such a ring is controlled by the rate at which it is radiatively heated, with a near balance prevailing between the accelerations produced by the Coriolis and magnetic tension forces in the $r$- and $\\theta$-directions.  Strongly magnetized rings with smaller cross-sectional radii experience larger initial buoyant accelerations and have shorter radiative heating time scales. As a result, they attain higher rise speeds and require less time to traverse the envelope than do rings with weaker fields and/or larger cross-sectional radii (see also Moss 1989).  A majority of the ring models enumerated in Table 1 reach the stellar surface in a time that is less than the main sequence lifetime of the 9 $M_\\odot$ star, the only exceptions being solutions with large values of $\\beta_0$ and/or $a_0$.  If the field strengths and flux tube sizes considered herein are realistic, this result suggests that the field produced by a core dynamo could manifest itself in the surface layers of a hot star at a relatively early stage of main sequence evolution.  In light of the discussion of \\S3, the strength of the field in a buoyant flux ring that reaches the photosphere of the 9 $M_\\odot$ stellar model is estimated to be $B_* \\approx B_0\\ (\\rho_{e*} / \\rho_{e0}) \\approx 9.44\\ (\\beta_0 / 10^4)^{-1/2}$ G.  The preceding conclusions are drawn from results obtained using a simple model for a thin, isolated, untwisted flux ring immersed in a non-evolving, rigidly rotating stellar interior.  Given the exploratory character of this investigation, a more detailed treatment of flux tube structure and dynamics is neither warranted nor practical.  However, several potentially important effects that have been neglected in the present analysis can be included within the context of the basic model of \\S2.  Among these are large-scale internal circulatory flows, the development (over time) of gradients in composition and angular velocity, the influence of radiation pressure on tubes inside more massive stars, mass loss, and the consequences of magnetic interference with radiative energy transport (e.g., thermal shadows and small-scale flows; Parker 1984).  Of the effects noted above, rotationally driven, meridional circulation may have the most significant impact on the results described in \\S3.  As noted in that section, each of the ring models listed in Table 1 spends considerably more time traveling through the thin shell $0.90\\ R_* \\leq r \\leq 0.96\\ R_*$ than in ascending from the starting radius to $r = 0.90\\ R_*$ (see Figure 9).  In solution 1, for example, the former displacement requires approximately $1.29 \\times 10^5$ years, while the latter occurs in only $1.38 \\times 10^3$ years, a difference in travel time of nearly a factor of 100.  We point out that the inclusion of meridional circulation may decrease the time needed to transport magnetic flux through the layers just below the photosphere.  Estimates of the circulation speed inside upper main sequence stars suggest that while such flows proceed quite slowly at great depth, they can become fast in the lower-density region near the surface (see, e.g., Kippenhahn \\& Weigert 1994).  Hence, flux tube advection by a circulatory flow may act to supplement buoyant transport in the outermost portion of the stellar interior, thereby shortening the time between the production of fields in the core and their emergence in the atmosphere.  The combined effects of buoyancy and advection will be investigated in the next paper of this series."
    },
    "0212/astro-ph0212230_arXiv.txt": {
        "abstract": "{We have investigated whether the $\\nu_4$ feature of $\\rm NH_4^+$ is a viable candidate for the 6.85 $\\mu$m absorption band seen towards embedded young stellar objects. To produce $\\rm NH_4^+$ astrophysical ice analogs consisting of $\\rm H_2O$, $\\rm CO_2$, $\\rm NH_3$ and $\\rm O_2$ were UV photolysed. The IR spectra reveal peaks that are identified with the NH$_4^+$, NO$_2^-$, NO$_3^-$ and HCO$_3^-$ ions. It is shown that the $\\rm NH_4^+$ matches two absorption features that are observed towards embedded young stellar objects, i.e., the strong 6.85 $\\mu$m feature and the 3.26 $\\mu$m feature.  The characteristic redshift with temperature of the interstellar 6.85 $\\mu$m feature is well reproduced. The abundance of $\\rm NH_4^+$ in interstellar ices would be typically 10 \\% relative to $\\rm H_2O$.  The experiments show that the counterions produce little distinct spectral signature but rather a pseudo-continuum if a variety of them is present in a $\\rm H_2O$ dominated environment. The anions could therefore go undetected in IR spectra of interstellar ice. In the ISM, where additional mechanisms such as surface chemistry and additional elements such as sulfur are available many acids and an even wider variety of anions could be produced.  These components may be detectable once the ices sublime, e.g, in hot cores. ",
        "introduction": "The nature of the 6.85 $\\mu$m absorption feature towards embedded Young Stellar Objects (YSO's) has remained an enigma since its discovery 25 years ago (Russell et al. 1977). While the first high resolution observations of this band were recently obtained by the Short Wavelength Spectrometer (SWS) on board the Infrared Space Observatory (ISO; Schutte et al. 1996, Dartois et al. 1999a; Keane et al. 2001), this did not yet help to clarify its origin. A number of candidates, such as carbonates, and the CH deformation modes in organic molecules like methanol, could be excluded based on the feature's spectral properties and the absence of strong additional bands. One possibility, however, the $\\nu_4$ mode of the ammonium ion ($\\rm NH_4^+$), first proposed by Grim et al. (1989b), could not be excluded. $\\rm NH_4^+$ is produced in astrophysical ice analogs by acid-base reactions.  This is achieved either by deposition and warm-up of $\\rm NH_3$ together with acids such as HNCO or HCOOH (Novozamsky et al. 2001, Schutte et al. 1999), or by photolysis of ice mixtures containing $\\rm NH_3$ (e.g., $\\rm NH_3$/$\\rm O_2$; Grim et al. 1989a) where the acids are produced by the photochemistry. Recently, the good spectral correspondence between the $\\nu_4$ feature of $\\rm NH_4^+$ and the interstellar 6.85 $\\mu$m absorption as seen by ISO/SWS was again demonstrated (Demyk et al. 1998). Nevertheless, this assignment faces a fundamental problem, since it requires a large abundance of $\\rm NH_4^+$ of $\\sim$ 10 \\% in the ices near YSO's. An equal amount of negative charge would need to be present. Although the ions $\\rm OCN^-$ (also referred to as NCO- in the chemistry literature) and (probably) $\\rm HCOO^-$ have been identified, their abundance falls far short of what is needed for the balance (Schutte et al. 1996a; Gibb et al. 2000; Keane et al. 2001; 2002). No sign of other negative solid state species has shown up in the ISO data.  This paper revisits the possibility that $\\rm NH_4^+$ is responsible for the interstellar 6.85 $\\mu$m band. We investigate in the laboratory whether a {\\it variety} of anions can supply the required countercharge, and that such a combination of species is either lacking in outstanding IR features, or fall below current detection limits. Taking the cosmic abundances of the elements into consideration, the counterions would have to consist of H, O, C, N and S in order to make a significant contribution. If we furthermore constrain the acid precursors to species with at most 1 carbon atom, possibilities would still include a variety of thermodynamically stable species, i.e., OH$^-$, CN$^-$, OCN$^-$, HCOO$^-$, $\\rm HCO_3^-$, $\\rm CO_3^{2-}$, NO$_2^-$, NO$_3^-$, $\\rm SO_4^{2-}$, HSO$_3^-$, and $\\rm SO_3^{2-}$. In principle an astrophysical ice analog containing $\\rm NH_4^+$ and a mixture of such negative ions could be prepared by depositing $\\rm NH_3$ together with a variety of acids. However, experimental limitations compel us to produce the ions {\\it in situ} with UV photolysis.  To this end, we processed mixtures of $\\rm H_2O$/$\\rm CO_2$/ $\\rm NH_3$/$\\rm O_2$ to produce acids such as $\\rm H_2CO_3$ (carbonic acid), $\\rm HNO_3$ (nitric acid) and $\\rm HNO_2$ (nitrous acid). These acids react with $\\rm NH_3$ forming $\\rm NH_4^+$ and a mixture of counterions, as desired. Earlier work on similar mixtures indeed showed the formation of such species (Moore \\& Khanna 1991; Gerakines et al. 2000; Grim et al. 1989b).  Of course, this experimental method of producing the ions may be different from what happens in the ISM. For example, the high abundance of $\\rm O_2$ in our samples, necessary to stimulate the photochemical production of oxygen rich acids, may considerably exceed the abundance in interstellar ices (Vandenbussche et al. 1999). In space acids are possibly produced by other mechanisms such as surface chemistry (e.g., Keane \\& Tielens 2002, in preparation).  To verify whether $\\rm NH_4^+$ is a plausible candidate for the identification of the 6.85 $\\mu$m band a number of crucial issues need to be adressed. First of all, detailed spectroscopy of the $\\nu_4$ mode is needed to investigate its behaviour as a function of ice composition and temperature. Next, the issue whether other infrared features of $\\rm NH_4^+$ could be present in interstellar ice spectra must be adressed. Subsequently, the counterions that are produced in the experiments should be identified to assess their astrophysical relevance. Finally, the spectral signature of the counterions under various conditions is studied, where the prime question is whether conditions exist at which their infrared signatures become inconspicuous.  The paper is organized as follows. In Sect. 2 we review the experimental techniques. Sect. 3 summarizes the results. $\\rm NH_4^+$ is created by UV photolysis of astrophysical ice analogs. We describe the spectral properties of $\\rm NH_4^+$ under such conditions.  Furthermore we establish which counterions are formed in the experiments and study their spectroscopic properties as well. In sect. 4 the obtained $\\rm NH_4^+$ spectra are compared with observations of YSO's. Besides matching the 6.85 $\\mu$m band, it is furthermore shown that the interstellar 3.26 $\\mu$m feature closely corresponds to one of the $\\rm NH_4^+$ features. Also, using the experimental results, it is investigated whether the absence of features due to counterions in the observations can be reconciled with an assignment of the 6.85 $\\mu$m band to $\\rm NH_4^+$. In Sect. 5, the astrophysical implications of the $\\rm NH_4^+$ identification are discussed.  Sect. 6, finally, summarizes the conclusions of this paper. ",
        "conclusions": "Experiments involving UV photolysis of $\\rm H_2O$/$\\rm CO_2$/$\\rm NH_3$/$\\rm O_2$ ice mixtures efficiently produce ammonium ($\\rm NH_4^+$) together with probably $\\rm HCO_3^-$, $\\rm NO_3^-$, and $\\rm NO_2^-$. These ions are produced by acid-base reactions between the $\\rm NH_3$ and photochemically produced acids. The $\\nu_4$ mode of $\\rm NH_4^+$ gives a strong feature at 6.85 $\\mu$m.  Other $\\rm NH_4^+$ features fall at 3.27 and 3.50 $\\mu$m. Due to strong interaction of the ion with the matrix, the $\\nu_4$ $\\rm NH_4^+$ band shifts strongly redward with increasing temperature. The infrared signature of the counterions is very weak for $\\rm H_2O$-dominated ices.  Comparison to the 6.85 $\\mu$m feature towards embedded young stellar objects shows that the $\\nu_4$ $\\rm NH_4^+$ feature obtained in our experiments provides a good match.  The observed red shift of the interstellar band with increasing line-of-sight dust temperature is also consistent with this assignment.  Furthermore, features are observed at 3.26 and 3.48 $\\mu$m towards these objects which could be (partly) due to the ammonium ion.  The implied abundance of $\\rm NH_4^+$ relative to $\\rm H_2O$ is 3 - 17 \\%.  The two anions observed in interstellar ices, $\\rm OCN^-$ and (possibly) $\\rm HCOO^-$, are insufficiently abundant to balance the positive charge from the $\\rm NH_4^+$ present.  Therefore heavier negative ions of the kind formed in our experiments appear to be required. It is yet unclear whether such ions are formed by energetic processing, as in the laboratory experiments, or by grain surface chemistry.  Future laboratory studies should include analysis by mass spectroscopy to exactly define the nature of the photochemically produced acids. Sulfur could be added to the mixture to extend the number of acids that are produced. On the observational side, the $\\rm NH_4^+$ assignment could be tested by probing the correlation between the 3.26 and 6.85 $\\mu$m bands.  In addition, it would be very interesting to search for rotational emission bands of evaporating acids like $\\rm HNO_2$, $\\rm H_2CO_3$ or $\\rm H_2SO_3$ in hot core regions."
    },
    "0212/astro-ph0212006_arXiv.txt": {
        "abstract": "Cosmic microwave background data shows the observable universe to be nearly flat, but leaves open the question of whether it is simply or multiply connected.  Several authors have investigated whether the topology of a multiply connect hyperbolic universe would be detectable when $0.9 < \\Omega < 1$.  However, the possibility of detecting a given topology varies depending on the location of the observer within the space.  Recent studies have assumed the observer sits at a favorable location.  The present paper extends that work to consider observers at all points in the space, and (for given values of $\\Omega_m$ and $\\Omega_\\Lambda$ and a given topology) computes the probability that a randomly placed observer could detect the topology. The computations show that when $\\Omega = 0.98$ a randomly placed observer has a reasonable chance ($\\sim 50\\%$) of detecting a hyperbolic topology, but when $\\Omega = 0.99$ the chances are low ($< 10\\%$) and decrease still further as $\\Omega$ approaches one. ",
        "introduction": "Analysis of recent cosmic microwave background (CMB) data suggests an approximately flat universe with the total energy density parameter $\\Omega$ almost surely lying in the range $0.9 < \\Omega < 1.1$.\\cite{sievers,netterfield}  The near flatness of the observable universe does not preclude a multiconnected spatial topology, but may push the topology to a scale larger than the horizon radius, making it difficult or impossible to detect. Recent studies have examined the extent to which a nontrivial topology may or may not be observable in a locally spherical universe\\cite{wlu02} or a locally hyperbolic one.\\cite{gomero1,aurichsteiner,inoue3,gomero2} In a locally flat universe there is no {\\it a priori} relationship between the topology scale and the horizon scale, so a great deal of luck would be required for the two to coincide.  In their most recent study of multiconnected hyperbolic universes,\\cite{gomero2} Gomero, Rebou{\\c c}as, and Tavakol examine the seven smallest known hyperbolic topologies and find that for a set of cosmological parameters given by Bond {\\it et al.}\\cite{bond} with $\\Omega = 0.99$, five of the seven topologies would be potentially detectable using CMB methods, while for a set of parameters given by Jaffe {\\it et al.}\\cite{jaffe} with $\\Omega = 0.98$, all seven would be potentially detectable.  A topology is considered potentially detectable by an observer at a point $p$ in the space if the observer's horizon radius\\footnote{One may substitute the last scattering surface at $z \\simeq 1100$ for the absolute horizon at $z = \\infty$ with little effect on the horizon radius.} $r_{hor}$ exceeds the injectivity radius\\footnote{Given a point $p$ in a multiconnected space, the {\\it injectivity radius} at $p$ is defined to be the radius of the largest ball centered at $p$ whose interior does not overlap itself.  Equivalently, twice the injectivity radius is the length of the shortest topologically nontrivial path from $p$ to itself.} $r_{inj}(p)$ at $p$.  In a 3-torus and in many spherical topologies, the injectivity radius $r_{inj}$ is constant throughout the whole space, so if the topology is potentially detectable by an observer at point $p$, it is detectable by any other observer at any other point $q$ in the same space. In hyperbolic topologies, by contrast, the injectivity radius varies from point to point, so a hyperbolic topology might be detectable by an observer at point $p$ (where the injectivity radius $r_{inj}(p)$ is small), but undetectable by a different observer at some other point $q$ (where the injectivity radius $r_{inj}(q)$ is large).  In their most recent work,\\cite{gomero2} Gomero, Rebou{\\c c}as, and Tavakol consider the question of detectability from the most favorable point in the space, that is, from a point of minimal injectivity radius.  The present article extends their work by considering the detectability of the topology at arbitrary locations in the space. The variation of the injectivity radius across the space will be summarized in an {\\it injectivity profile} showing what fraction of the space's volume has a given injectivity radius.  Combining the injectivity profile (Section~\\ref{SectionInjectivityProfiles}) with an estimate for the horizon radius (Section~\\ref{SectionHorizonRadius}) reveals the probability that a randomly placed observer could potentially detect the topology, that is, it tells the fraction of the manifold's volume in which $r_{hor} > r_{inj}$ (Section~\\ref{SectionResults}). ",
        "conclusions": "In the case $\\Omega_{total} = 0.98$, the horizon radius $r_{hor} = 0.43$ is large enough that an observer at a random location in a small hyperbolic universe would have a reasonable chance of detecting the topology.  However, in the case $\\Omega_{total} = 0.99$, the horizon radius $r_{hor}$ drops to $0.30$ and a random observer would have little or no chance of detecting the topology.  There exist closed hyperbolic 3-manifolds with arbitrarily short closed geodesics, so no matter how close $\\Omega_{total}$ is to one, there will always be infinitely many hyperbolic manifolds in which well-placed observers could detect the topology. However, the probability that an observer could detect the topology from a random point in such a manifold decreases in proportion to $1 - \\Omega_{total}$."
    },
    "0212/astro-ph0212376_arXiv.txt": {
        "abstract": "The mass--energy formula of black holes implies that up to 50\\% of the energy can be extracted from a static black hole. Such a result is reexamined using the recently established analytic formulas for the collapse of a shell and expression for the irreducible mass of a static black hole. It is shown that the efficiency of energy extraction process during the formation of the black hole is linked in an essential way to the gravitational binding energy, the formation of the horizon and the reduction of the kinetic energy of implosion. Here a maximum efficiency of 50\\% in the extraction of the mass energy is shown to be generally attainable in the collapse of a spherically symmetric shell: surprisingly this result holds as well in the two limiting cases of the Schwarzschild and extreme Reissner-Nordstr\\\"{o}m space-times. Moreover, the analytic expression recently found for the implosion of a spherical shell onto an already formed black hole leads to a new exact analytic expression for the energy extraction which results in an efficiency strictly less than 100\\% for any physical implementable process. There appears to be no incompatibility between General Relativity and Thermodynamics at this classical level. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/astro-ph0212140_arXiv.txt": {
        "abstract": "Recent observations have revealed many new puzzles related to core-collapse supernovae, including the formation of magnetars and black holes and their possible GRB connections. We review our current understanding of the origin of pulsar kicks and supernova asymmetry. It is argued that neutron star kicks are intimately connected to the other fundamental parameters of young neutron stars, such as the initial spin and magnetic field strength. ",
        "introduction": "The subject of supernovae (SNe) has a long history, but the modern era of SN research really began in 1934 when Baade and Zwicky made the prophetic suggestion that the death of massive stars, SN explosions and neutron star (NS) formation are connected events. This suggestion was confirmed by the discovery of the pulsar in the Crab supernova remnant (SNR) in 1968; the SN, as is well known, was actually observed in 1054 by Chinese astronomers.  Today the mechanism of SN explosion remains an unsolved problem. Moreover, observations over the last few years suggest that we may actually know less than we thought about core collapse and explosion of massive stars. Here we discuss a small sample of unsolved problems related to SNe and NS formation, focusing on the problem of NS kicks.  \\smallskip {\\it Basic Paradigm for Core-Collapse Supernovae:} The current paradigm for core-collapse supernovae is that they are neutrino-driven (see, e.g., Bethe 1990; Janka et al.~2001; Burrows \\& Thompson 2002 for reviews): As the central core of a massive star collapses to nuclear density, it rebounds and sends off a shock wave, leaving behind a proto-NS. The shock stalls at several 100's km because of neutrino loss and nuclear dissociation in the shock. A fraction of the neutrinos emitted from the proto-NS get absorbed by nucleons behind the shock, thus reviving the shock, leading to an explosion on the timescale several 100's ms --- This is the so-called ``delayed mechanism''. However, 1D simulations with detailed neutrino transport seem to indicate that neutrino heating of the stalled shock, by itself, does not lead to a robust explosion (e.g., Rampp \\& Janka 2000; Liebendoerfer et al.~2002). It has been argued that neutrino-driven convection in the proto-NS (which tends to increase the neutrino flux) and that in the shocked mantle (which tends to increase the neutrino heating efficiency) are central to the explosion mechanism, although it is not clear how robust of these convections are (see, e.g., Mezzacappa et al.~1998; Fryer \\& Warren 2002). Clearly, in this ``standard model'', the problem of SN explosion is a quantitative one, involving 3D radiation (neutrino) (and possibly relativistic) hydrodynamics.  This basic picture of core-collapse SNe, however, is likely to be incomplete. There are some obvious puzzles as a result of recent observations:  \\smallskip {\\it Formation of Magnetars:} Observations over the last few years have revealed two new classes of young NSs, the soft gamma-ray repeaters (SGRs) and anomalous X-ray pulsars (AXPs). They have many common characteristics and are mostly likely related to each other (see Kaspi's contribution in this proceedings). Strong physical arguments suggest that these are young NSs endowed with superstrong magnetic fields $B\\go 10^{14}$~G (see Thompson 2001). The existence of these magnetars poses some obvious questions: Under what conditions core-collapse of massive stars will lead to radio pulsars vs. magnetars? What is the branching ratio? What is the origin of the NS magnetic field? Does the B-field (in combination with rotation) play any dynamical role in the explosion? Currently we do not have firm answers to these questions (see, e.g., Thompson \\& Duncan 1993).  \\smallskip {\\it Black Hole Formation and Core-Collapse SNe:} BHs are also formed in the core collapse of massive stars. Recent observations showed that BH formation can be accompanied by SN explosion, at least in two cases: The companion of the BH X-ray binary GRO J1655-40 (Nova Sco) and that of SAX J1819.3-2525 (V4641 SGR) have high abundance of $\\alpha$-elements (Israelian et al.~1999; Orosz et al.~2001); these $\\alpha$-elements can only be produced in a SN explosion. Apparently, the companion stars have been polluted by material ejected in the SN that accompanied the formation of the BH primary (see Podsiadlowski et al.~2002). These observations pose a host of questions: What are the differences between a SN that made a NS and a SN that made a BH? How is the BH formed? We could have a direct collapse to BH in which the shock wave never successfully makes it to induce an explosion, or we could have a indirect process where a shock wave successfully makes an explosion and a NS forms temporarily followed by fall-back, or loss of angular momentum and thermal energy in the proto-NS which then collapses to a BH. This indirect process may explain the the relatively large space velocity of GRO J1655-40.  On a more speculative side, there is possible connection between SN and the central engine of gamma-ray bursts (GRBs). In the last few years a growing list of observations suggests that some GRBs are connected with the death of massive star and SNe. The famous case is GRB 980425, which coincided both in time and in position with a Type Ic SN1998bw. In at least 3 GRBs (980326,~970228,~000911), the rebrightening of the optical afterglows about a month after the initial bursts have been observed and interpreted as the underlying SNe which emerged when the afterglows faded. Possible emission line features (K$_\\alpha$ of Fe, O, Mg, Si, etc) in several X-ray afterglows at hours-days may also be an indication of SNe. The obvious question is: under what conditions will the collapse of a massive star lead to GRB, SN with BH, SN with NS, etc.? ",
        "conclusions": ""
    },
    "0212/hep-th0212027_arXiv.txt": {
        "abstract": "It has recently been proposed in Refs.~\\cite{MBK01,BFM02,BM02} that the dark energy could be attributed to the cosmological properties of a scalar field with a non-standard dispersion relation that decreases exponentially at wave-numbers larger than Planck scale ($k_{\\rm phys}>M_{\\rm Pl}$). In this scenario, the energy density stored in the modes of trans-Planckian wave-numbers but sub-Hubble frequencies produced by amplification of the vacuum quantum fluctuations would account naturally for the dark energy. The present article examines this model in detail and shows step by step that it does not work. In particular, we show that this model cannot make definite predictions since there is no well-defined vacuum state in the region of wave-numbers considered, hence the initial data cannot be specified unambiguously. We also show that for most choices of initial data this scenario implies the production of a large amount of energy density (of order $\\sim M_{\\rm Pl}^4$) for modes with momenta $\\sim M_{{\\rm Pl}}$, far in excess of the background energy density. We evaluate the amount of fine-tuning in the initial data necessary to avoid this back-reaction problem and find it is of order $H/M_{{\\rm Pl}}$. We also argue that the equation of state of the trans-Planckian modes is not vacuum-like. Therefore this model does not provide a suitable explanation for the dark energy. ",
        "introduction": "\\label{Sec:I}  Recently, it has been claimed in a series of articles~\\cite{MBK01,BFM02,BM02} that the cosmic dark energy component could be explained naturally by the trans-Planckian energy of a scalar field with a suitable non-linear dispersion relation in the trans-Planckian regime. Such dispersion relations, which relate the frequency $\\omega_{\\rm phys}$ to the wave-number $k_{\\rm phys}$ of a scalar field wave-packet, and which depart from the standard linear dispersion relation in the trans-Planckian regime are a way of modeling phenomenologically the unknown physics for sub-Planckian wavelengths. They have been used extensively in the recent literature in the context of black-hole physics~\\cite{U95} and of the inflationary trans-Planckian problem~\\cite{MB}.  In the particular case considered in Refs.~\\cite{MBK01,BFM02,BM02}, the dispersion relation departs from its standard linear form and approach a decreasing exponential at large wave-numbers. This type of dispersion relation could possibly emerge from string theory~\\cite{BFM02}. It has been argued that the energy density of the modes of sub-Planckian wavelengths and sub-Hubble frequencies (referred to as ``tail'' modes) is naturally of the same order as the critical energy density today and has the same equation of state as a cosmological constant.  Hence, it could account without fine-tuning for the dark energy.  The energy density contained in the tail today has been calculated in Ref.~\\cite{MBK01} by solving for the time evolution of a test quantum scalar field evolving in the curved cosmological background, assuming its initial state at the onset of inflation is the vacuum. The equation of state of the tail modes has been calculated in Ref.~\\cite{BM02} and its cosmological evolution has been solved to argue that the cosmic coincidence problem (why does the dark energy dominate now?) is solved.  In this paper we argue that this model does not and can not work for several reasons. We first argue that the energy density and equation of state of the ``tail'' modes depend directly on the choice of initial data for the scalar field, and that this latter cannot be specified {\\it unambiguously} since there is no preferred initial vacuum state (Section~\\ref{Sec:II}). This implies that any cosmological consequence derived depends directly on the ad-hoc choice of initial data (initial quantum state). We then show that the violation of the Wentzel-Kramers-Brillouin (WKB) approximation in the remote past for all co-moving wave-numbers, which is inevitable in the present scenario, implies the continuous production {\\it at all times} of a large amount of quanta with physical wave-number $\\sim M_{{\\rm Pl}}$ (Section~\\ref{Sec:III}). This finding is in agreement with general arguments given by Starobinsky~\\cite{S01} (this latter work did not however study the present scenario). We evaluate the amount of energy density produced for modes of physical wave-number and frequency $\\sim M_{{\\rm Pl}}$ for various choices of initial data and conclude that it is generically of order $M_{{\\rm Pl}}^4$. This process takes place at all times, and since the energy density produced is much larger than the background energy density, it implies that the semi-classical perturbative framework on which the model of Refs.~\\cite{MBK01,BFM02,BM02} rests breaks down.  In an earlier study, devoted to constructing an effective stress-energy tensor for theories with non-linear dispersion relations~\\cite{LMMU01}, we already criticized this model by arguing that it led generically to the wrong equation of state. We revisit this issue further in Section~\\ref{Sec:IV}, where we prove that the effective energy-momentum tensor we derived earlier is well-behaved, thus disproving an improper claim of Ref.~\\cite{BM02} and confirming our earlier criticisms. We provide a summary of our conclusions in Section~\\ref{Sec:V}. ",
        "conclusions": "\\label{Sec:V}  In this article, we have examined in detail the scenario proposed in Refs.~\\cite{MBK01,BFM02,BM02} which attributes the dark energy to the properties of a scalar field with a dispersion relation that decreases exponentially with trans-Planckian wave-number $k_{\\rm phys}\\gtrsim k_{\\rm c}\\sim M_{\\rm Pl}$. We have demonstrated that this mechanism does not work, mainly for two reasons. (i) The mode function of the scalar field does not behave as a plane wave as $\\eta\\rightarrow-\\infty$ in the so-called ``tail'' (i.e., the part of the non-linear dispersion relation where $\\omega_{\\rm phys}\\ll H$ and $k_{\\rm phys}\\gg M_{{\\rm Pl}}$). This implies that there is no definite prescription for constructing a well-defined initial vacuum state, hence that the choice of initial data is entirely arbitrary. This situation is similar to the problem of setting initial data for cosmological perturbations in the absence of an accelerated expansion era (but with linear dispersion relations), in which case the data would have to be specified on super-horizon scales where the mode function does not oscillate and is frozen by the expansion. In the scenario of Refs.~\\cite{MBK01,BFM02,BM02} the WKB approximation is not valid at all times for modes in the tail. This explains that the notion of adiabatic vacuum cannot be used to set up initial data and the initial state proposed in Ref.~\\cite{MBK01} is thus ad-hoc. It also brings us to the second objection against this scenario: (ii) since all modes originate in the tail of the dispersion relation, the breakdown of the WKB approximation for a given physical wave-number at early times implies the continuous production of a substantial number of quanta with physical wave-numbers $\\gtrsim k_{\\rm c}$. The breakdown of the WKB approximation can indeed be seen as the signal of a strongly non-adiabatic evolution which is generically associated with particle production in expanding space-times. We have calculated the amount of energy density produced in modes of physical wave-number $k_{\\rm c}$ as a function of the two parameters that characterize the (arbitrary) choice of initial data. We have shown that this energy density is generically of order $M_{{\\rm Pl}}^4$.  The production of energy density in quanta with wave-numbers $\\sim M_{{\\rm Pl}}$ in excess of the background energy density $\\sim H^2 M_{{\\rm Pl}}^2$ implies the breakdown of the perturbative semi-classical framework used for the calculation, and renders all claims irrelevant.  There is a small region of parameter space in which this energy density can be tuned down to zero, but the fine-tuning in the choice of initial data is of order $H/M_{{\\rm Pl}}$. Such a fine-tuning at the time of inflation is of order $\\sim 10^{-6}$ and is probably acceptable, but today, it is of the same order as the celebrated fine-tuning of the cosmological constant problem."
    },
    "0212/astro-ph0212410_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}} \\newcommand{\\TabCapp}[2]{\\begin{center}\\parbox[t]{#1}{\\centerline{ \\small {\\spaceskip 2pt plus 1pt minus 1pt T a b l e} \\refstepcounter{table}\\thetable} \\vskip2mm \\centerline{\\footnotesize #2}} \\vskip3mm \\end{center}}  \\newcommand{\\MakeTableSepp}[4]{\\begin{table}[p]\\TabCapp{#2}{#3}\\vspace*{-.7cm} \\begin{center} \\TableFont \\begin{tabular}{#1} #4 \\end{tabular}\\end{center}\\end{table}}  \\newfont{\\bb}{ptmbi8t at 12pt} \\newfont{\\bbb}{cmbxti10} \\newfont{\\bbbb}{cmbxti10 at 9pt} \\newcommand{\\uprule}{\\rule{0pt}{2.5ex}} \\newcommand{\\douprule}{\\rule[-2ex]{0pt}{4.5ex}} \\newcommand{\\dorule}{\\rule[-2ex]{0pt}{2ex}} \\def\\thefootnote{\\fnsymbol{footnote}}  \\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 the catalog of RR~Lyrae stars from 2.4 square degrees of central parts of the Small Magellanic Cloud (SMC). The photometric data were collected during four years of the OGLE-II microlensing survey. Photometry of each star was obtained using the Difference Image Analysis (DIA) method. The catalog contains 571 objects, including 458~RRab, 56~RRc variables, and 57 double mode RR~Lyr stars (RRd). Additionally we attach a~list of variables with periods between 0.18--0.26 days, which are probably $\\delta$~Sct stars. Period, {\\it BVI} photometry, astrometry, amplitude, and parameters of the Fourier decomposition of the {\\it I}-band light curve are provided for each object. We also present the Petersen diagram for double mode pulsators.  We found that the SMC RR~Lyr stars are fairly uniformly distributed over the studied area, with no clear correlation to any location. The most preferred periods for RRab and RRc stars are 0.589 and 0.357 days, respectively. We noticed significant excess of stars with periods of about 0.29 days, which might be second-overtone RR~Lyr stars (RRe). The mean extinction free magnitudes derived for RRab stars are 18.97, 19.45 and 19.73 mag for the $I$, $V$ and {\\it B}-band, respectively.  All presented data, including individual {\\it BVI} observations, are available from the OGLE {\\sc Internet} archive.}{} ",
        "introduction": "\\vskip13pt The Magellanic Clouds provide an ideal opportunity to study in detail the structure and evolution of stars. Rich populations of stars approximately at the same, relatively small distance with small interstellar reddening make the Large and Small Magellanic Clouds very important targets for observing surveys.  RR~Lyr stars were first discovered in the SMC near the cluster NGC121 by Thackeray (1951). The stars were more than a~magnitude fainter than expected, what was a~crucial confirmation of the major revision of the extragalactic distance scale proposed by Baade (1952). The first systematic search for field RR~Lyr variables in the SMC was conducted by Graham (1975). During that survey 76 RR~Lyr stars were discovered in a~${1\\arcd\\times1\\zdot\\arcd3}$ outlying field centered on the cluster NGC121. Smith \\etal (1992) used {\\it B}-band photographic photometry of a~field in the northeast arm of the SMC to identify additional 22 probable RR~Lyr variables. The same outlying region of the SMC was observed by Sharpee \\etal (2002). They presented {\\it V}-band and {\\it B}-band CCD photometry of a~few RR~Lyr stars.  Walker (1989) surveyed five SMC clusters for RR~Lyr stars. Four already known RR~Lyr variables in the NGC121 were rediscovered, but no such stars were found in the other clusters. Because the age of NGC121 was estimated to be $12\\pm2$~Gyr (Stryker \\etal 1985), and Lindsay~1, the next oldest SMC cluster, $10\\pm2$~Gyr (Olszewski \\etal 1987), Walker concluded that the minimum age of RR~Lyr stars is about 11~Gyr.  In 1990s the number of known variable stars in the Magellanic Clouds dramatically increased, when the large microlensing searches began regular photometric monitoring of both galaxies (\\eg MACHO -- Alcock \\etal 1993, EROS -- Aubourg \\etal 1993). Natural by-product of the microlensing surveys are huge databases with precise photometry of millions of stars.  The SMC was also included to the list of targets of the second phase of the Optical Gravitational Lensing Experiment (OGLE-II; Udalski, Kubiak and Szyma{\\'n}ski 1997). About 2.4 square degrees of central part of the SMC were observed each night during the observing seasons 1997--2000. Photometry was obtained with the $BVI$ filters, closely resembling the standard system.  The OGLE-II survey has yielded a~particularly rich harvest of variable stars from the SMC. In the previous papers we presented the catalog of eclipsing binary stars (Udalski \\etal 1998b), the catalog of Cepheids from the SMC (Udalski \\etal 1999), and general catalog of variable stars detected in the Magellanic Clouds ({\\.Z}ebru{\\'n} \\etal 2001). In addition {\\it BVI} maps of the SMC were released providing precise photometry and astrometry of about 2.2 million stars (Udalski \\etal 1998a).  In this paper we present a sample of 571 RR~Lyr stars and several other pulsating objects, likely $\\delta$~Sct stars, detected in the OGLE-II fields in the SMC. The stars were selected from the reprocessed OGLE-II photometry based on the Difference Image Analysis (DIA) technique -- Wo{\\'z}niak's (2000) implementation of Alard and Lupton (1998) and Alard (2000) optimal Point Spread Function matching algorithm.  Similarly to the previous catalogs, all data presented in this paper, including individual observations, are available to the astronomical community from the OGLE {\\sc Internet} archive. ",
        "conclusions": "The long term photometric surveys of dense stellar systems may provide large samples of all types of variable stars (Paczy{\\'n}ski 2000). In this paper we present the catalog of RR~Lyr variable stars -- objects playing a major role in studies of the processes of stellar pulsation, post-main-sequence evolution of stars, calibration of extragalactic distances, and structure and age of galaxies. This is the largest sample of RR~Lyr variables detected in the SMC so far.  We reported detection of 57 double-mode RR~Lyr stars. We also found RR~Lyr variables exhibiting two closely spaced frequencies, most probably related to non-radial pulsations.  About 10\\% of the detected RR Lyr variables are RRc stars. Next 10\\% of the \\renewcommand{\\TableFont}{\\footnotesize} \\setcounter{table}{8} \\MakeTable{ l@{\\hspace{3pt}} l@{\\hspace{6pt}} l@{\\hspace{6pt}} l@{\\hspace{4pt}} c@{\\hspace{4pt}} c@{\\hspace{4pt}} }{12.5cm}{RR Lyr stars close to the SMC clusters} {\\hline \\noalign{\\vskip3pt} \\multicolumn{1}{c}{Field} & \\multicolumn{1}{c}{Star ID} & Cluster & Alternative & Cluster & Distance from \\\\ & & name & cluster & radius & the center \\\\ & & OGLE-CL- & name & [\\arcs] & [\\arcs]  \\\\ \\noalign{\\vskip3pt} \\hline \\noalign{\\vskip3pt} SMC$\\_$SC2 & OGLE004222.96$-$733223.8 & SMC0014 & & ~24 & 33 \\\\ SMC$\\_$SC3 & OGLE004504.31$-$725623.9 & SMC0024 & & ~57 & 68 \\\\ SMC$\\_$SC4 & OGLE004526.76$-$732801.7 & SMC0027 & B39 & ~36 & 51 \\\\ SMC$\\_$SC4 & OGLE004610.66$-$730355.6 & SMC0183 & & ~10 &  ~~1 \\\\ SMC$\\_$SC4 & OGLE004817.88$-$731815.2 & SMC0048 & B47 & ~49 & 67 \\\\ SMC$\\_$SC5 & OGLE004833.02$-$731800.0 & SMC0048 & B47 & ~49 & 25 \\\\ SMC$\\_$SC5 & OGLE004907.44$-$730617.5 & SMC0052 & H86-104 & ~18 & 25 \\\\ SMC$\\_$SC5 & OGLE004913.17$-$730651.0 & SMC0052 & H86-104 & ~18 & 21 \\\\ SMC$\\_$SC5 & OGLE004922.63$-$731220.9 & SMC0053 & & ~30 & 30 \\\\ SMC$\\_$SC5 & OGLE004924.05$-$732231.3 & SMC0054 & L39 & ~27 & 30 \\\\ SMC$\\_$SC5 & OGLE004924.31$-$731447.2 & SMC0195 & BS42 & ~28 & 35 \\\\ SMC$\\_$SC5 & OGLE004954.79$-$730321.3 & SMC0057 & B52 & ~49 & 67 \\\\ SMC$\\_$SC5 & OGLE004936.31$-$725229.2 & SMC0058 & H86-109 & ~36 & 51 \\\\ SMC$\\_$SC5 & OGLE005120.92$-$730920.1 & SMC0069 & NGC290 & ~36 & 36 \\\\ SMC$\\_$SC6 & OGLE005202.85$-$725707.3 & SMC0074 & K29,L44 & ~42 & 44 \\\\ SMC$\\_$SC6 & OGLE005243.99$-$724740.5 & SMC0085 & B66 & ~25 & 17 \\\\ SMC$\\_$SC7 & OGLE005549.04$-$725335.1 & SMC0105 & H86-165 & ~44 & 54 \\\\ SMC$\\_$SC7 & OGLE005612.31$-$731222.5 & SMC0106 & & ~26 & 15 \\\\ SMC$\\_$SC8 & OGLE005820.37$-$723841.4 & SMC0113 & L61-331 & ~24 & 19 \\\\ SMC$\\_$SC11 & OGLE010728.50$-$724527.6 & SMC0156 & L80 & ~41 & 42 \\\\ SMC$\\_$SC11 & OGLE010836.97$-$725321.6 & SMC0159 & NGC419 &102 & 80 \\\\ \\noalign{\\vskip3pt} \\hline} whole sample are RRd variables. Relatively small number of the first-overtone pulsators and the value of the average period of RRab stars (0.589 days) places the SMC RR~Lyr stars near the Oosterhoff type~I group.  Additionally, we conducted a search of RR~Lyr stars in the SMC clusters. Discovery of RR~Lyr variables in the clusters younger than NGC121 (hosting four RR~Lyr stars), would have a significant consequences on our understanding of evolution of old, solar-mass stars. We used the OGLE catalog of the SMC clusters (Pietrzy{\\'n}ski \\etal 1998) containing 238 clusters, their coordinates and angular sizes. We scanned our sample of RR~Lyr variables looking for stars in a distance smaller than 1.5 cluster radius from the cluster center. We obtained the list of 21 RR~Lyr stars located close to the line-of-sight to, in total, 19 SMC clusters. Only in two cases (OGLE-CL-SMC0048 and OGLE-CL-SMC0052) we detected two RR~Lyr stars in the cluster neighborhood. It is very probable that the majority of these cases are just  optical coincidences, because more or less similar number of stars could be expected in the background of SMC clusters (comparing the area of the clusters to the area of all our fields). Nevertheless, because of potential importance of the discovery of RR~Lyr variables in the SMC clusters, we provide in Table~9 the list of stars located close to the SMC clusters.  \\Acknow{We would like to thank Prof.~Bohdan Paczy{\\'n}ski for many discussions and valuable suggestions. We are very grateful to Prof.~Wojciech Dziembowski for helping us in the classification of stars. The paper was partly supported by the Polish KBN grant BST to Warsaw University Observatory. Partial support for the OGLE   project was provided with the NSF grants AST-9820314 and AST-0204908 and NASA grant NAG5-12212 to B.~Paczy\\'nski. We acknowledge usage of the Digitized Sky Survey which was produced at the Space Telescope Science Institute based on photographic data obtained using the UK Schmidt Telescope, operated by the Royal Observatory Edinburgh.}"
    },
    "0212/astro-ph0212283_arXiv.txt": {
        "abstract": "How massive were the first stars? This question is of fundamental importance for galaxy formation and cosmic reionization.  Here we consider how protostellar feedback can limit the mass of a forming star. For this we must understand the rate at which primordial protostars accrete, how they and their feedback output evolve, and how this feedback interacts with the infalling matter. We describe the accretion rate with an ``isentropic accretion'' model: $\\mds$ is initially very large ($0.03\\smyr$ when $m_*=1\\sm$) and declines as $m_*^{-3/7}$. Protostellar evolution is treated with a model that tracks the total energy of the star. A key difference compared to previous studies is allowance for rotation of the infalling envelope. This leads to photospheric conditions at the star and dramatic differences in the feedback. Two feedback mechanisms are considered: HII region breakout and radiation pressure from Lyman-$\\alpha$ and FUV photons. Radiation pressure appears to be the dominant mechanism for suppressing infall, becoming dynamically important around 20~$\\sm$. ",
        "introduction": "Recent numerical studies have followed the gravitational collapse of perturbations from cosmological to almost stellar dimensions\\cite{Abel:2002,Bromm:1999}.  Baryon-dominated clouds, cooled to about 200-300~K by trace amounts of molecular hydrogen, form at the centers of dark matter halos. For $n_{\\rm H}>10^4\\:{\\rm cm^{-3}}$ the cooling rate becomes independent of density, and so the dissipation of gravitational energy in the densest regions gradually raises the temperature.  In the simulation of Abel et al.\\cite{Abel:2002}, the gas cloud is centrally-concentrated ($\\rho\\propto r^{-k_\\rho}; \\: k_\\rho\\simeq2.2$) and is contracting quasi-hydrostatically with infall speeds about one third of the sound speed. In addition to thermal support, the cloud is filled with a turbulent cascade of weak shocks (T. Abel, private comm.).  The structure can be described by an approximately hydrostatic polytrope with $\\gamma_p=1+1/n=k_P/k_\\rho=2(1-1/k_\\rho)=1.1$, where $P\\propto r^{-k_P}$.  The contraction is akin to the maximally sub-sonic Hunter\\cite{Hunter:1977} settling solution, for which the accretion rate is a factor $\\phi_*\\simeq 2.6$ greater than the classic Shu\\cite{Shu:1977} solution. This accretion rate can be expressed in terms of the entropy parameter of the polytrope, $K=P/\\rho^{\\gamma_p}$ and the collapsed mass, $M\\simeq m_*$, \\beq \\mds=\\frac{8\\phi_*}{\\surd 3}\\left[\\frac{(3-\\krho)k_P^3 K^3}{2(2\\pi) ^{5-3\\gamma_p}G^{3\\gamma_p-1}}\\right]^{\\frac{1} {2(4-3\\gamma_p)}} M^{j}\\rightarrow0.026 K'^{15/7} \\left(\\frac{M}{M_\\odot}\\right)^{-3/7}\\smyr, \\label{eq:mds3} \\eeq where $j\\equiv 3(1-\\gamma_p)/(4-3\\gamma_p)$. The numerical evaluation assumes $\\phi_*=2.6$, since $\\gamma_p$ is not too different from one, and $K=1.88\\times10^{12}(T/300\\:{\\rm K}) n_{\\rm H,4}^{-0.1}\\:{\\rm cgs}\\equiv1.88\\times10^{12}K^\\prime\\:{\\rm cgs}$. We set the temperature normalization a factor $4/3$ higher than is seen in simulations\\cite{Abel:2002}, to allow for partial pressure support from sonic and isotropic turbulent motions; i.e. $T$ is an effective temperature. The small ratio of turbulent to thermal support is in marked contrast to contemporary massive star formation\\cite{McKee:2002}. In primordial clouds it is the microphysics of $\\rm H_2$ cooling that determines both the evolution (via $\\gamma_p$) and normalization (via $K^\\prime$ and $\\phi_*$) of the accretion rate.  Collapse assuming constant $K$ - ``isentropic accretion'' - agrees with 1-D numerical studies\\cite{Omukai:1998,Ripamonti:2002} (Fig. 1a).  The mass-averaged rotation speed is a fraction, $f_{\\rm Kep}\\simeq 0.5$, of Keplerian, approximately independent of radius\\cite{Abel:2002}. Assuming angular momentum is conserved inside the sonic point, $r_{\\rm sp}$, leads to a disk size $r_d = f_{\\rm Kep}^2 r_{\\rm sp}=3.4 (f_{\\rm Kep}/0.5)^2 (M/M_\\odot)^{9/7} K'^{-10/7}\\:{\\rm AU}$. Matter falls onto this disk at all radii $r\\leq r_d$, as well as directly to the star. We follow Ulrich\\cite{Ulrich:1976} in describing the density distribution of the rotating, freely-falling envelope.  \\begin{figure} \\includegraphics[height=.8\\textheight]{mary} \\caption{ (a) Protostellar accretion rate as a function of the collapsed mass ($\\simeq m_*$ in these models). {\\it Solid} line: fiducial isentropic accretion model ($K^\\prime=1$) from eq. (\\ref{eq:mds3}); dotted\\cite{Omukai:1998} and dashed lines extrapolated from 1-D numerical studies; long-dashed line is $\\mds=4.4\\times10^{-3}\\smyr$ used in the protostellar evolution models of Stahler et al.\\cite{Stahler:1986} and Omukai \\& Palla\\cite{Omukai:2001}. (b) Evolution of protostellar radius (lower, thick lines), which is the location of the accretion shock, and photospheric radius (upper, thin lines). The spherical constant accretion rate test case (dotted) is compared to other calculations\\cite{Stahler:1986,Omukai:2001} of $r_*$ (squares) for the evolution before H burning. Also shown are the spherical (dashed lines) and rotating ($f_{\\rm Kep}=0.5$, solid lines) isentropic accretion models. The initial condition is taken from 1-D hydrodynamical collapse simulations\\cite{Ripamonti:2002}. Note that the photospheric and protostellar radii are the same in the rotating model for $m_*>0.3\\sm$.(c) Geometry of the HII region (shaded) at polar breakout when $r_{\\rm HII}=r_g$, the gravitational radius for the ionized gas sound speed.  The protostar is at (0,0) and the disk is in the $z=0$ plane. Dashed streamlines show infall. (d) Mass scales of feedback processes versus the rotation parameter, $f_{\\rm Kep}$. HII region breakout at the pole and just above the equator occur at masses traced by the lower and upper dashed lines, respectively. Ly-$\\alpha$ and FUV radiation pressure becomes greater than twice the radial infall ram pressure (evaluated at angle $\\pi/3$ from the pole) for masses above the solid line.} \\end{figure} ",
        "conclusions": ""
    },
    "0212/astro-ph0212556_arXiv.txt": {
        "abstract": "We re-examine the Boussinesq hypothesis of an effective turbulent viscosity within the context of simple closure considerations for models of strong magnetohydrodynamic turbulence. Reynolds-stress and turbulent Maxwell-stress closure models will necessarily introduce a suite of transport coefficients, all of which are to some degree model-dependent. One of the most important coefficients is the relaxation time for the turbulent Maxwell stress, which until recently has been relatively ignored. We discuss this relaxation within the context of magnetohydrodynamic turbulence in steady high Reynolds-number, high magnetic-Reynolds-number shearing flows. The relaxation time for the turbulent Maxwell stress is not limited by the shear time scale, in contrast with Reynolds-stress closure models for purely hydrodynamical turbulence. The anisotropy of the turbulent stress tensor for magnetohydrodynamic turbulence, even for the case of zero mean-field considered here, can therefore not be neglected. This shear-generated anisotropy can be interpreted as being due to an effective turbulent elasticity, in analogy to the Boussinesq turbulent viscosity. We claim that this turbulent elasticity should be important for any astrophysical problem in which the turbulent stress in quasi-steady shear has been treated phenomenologically with an effective viscosity. ",
        "introduction": "The Boussinesq hypothesis of an effective turbulent viscosity is still widely used both in engineering applications and in astrophysics. This is true despite the passage of over one hundred years since the birth of the concept. There are a very large number of models for turbulent stress that give substantially better results than a simple effective viscosity, and in some situations, such as strongly time-dependent flows, the use of an effective viscosity gives notoriously bad results. Despite this, the concept of an effective turbulent viscosity has proven to be so persistent because it is useful, simple, and because it is based on an appealing analogy between the random motions of turbulent blobs and the kinetic theory of thermal motions of molecules. Thus, effective turbulent viscosity is commonly invoked as a good crude ``lowest-order'' approximation, for problems in which the Reynolds number is very large.  For problems in which the magnetic Reynolds number is also very large (as is often the case in astrophysics), we ask if a picture of turbulence as colliding blobs of fluid is at all useful even as a crude approximation. A more appropriate analogy for the turbulent component of the {\\em field}  may be not the kinetic theory of point-like molecules, but rather the kinetic theory of long polymers in melt or in solution \\citep{Ogil:2001, Will:2001}. For the crudest possible model, such an analogy would suggest the concept of an effective turbulent elasticity, in addition to the effective turbulent viscosity. Thus, corresponding to the Boussinesq approximation in hydrodynamic turbulence, the corresponding ``lowest-order'' approximation for the stress in MHD turbulence should include an additional coefficient, in analogy to a simple viscoelastic fluid. Just as the inclusion of elasticity in the stress response of an ordinary laboratory fluid can have dramatic consequences, we suggest that {\\em any} astrophysical problem in which turbulence has been treated phenomenologically with the Boussinesq viscosity might have dramatically different solutions if an effective elasticity were included as well.   In particular, for the purposes of quasi-steady shear, we claim that the anisotropy of the stress tensor is of much greater significance in the case of MHD turbulence than in the case of purely hydrodynamic turbulence. Thus, while investigations of the idealized problem of homogeneous istotropic hydrodynamic turbulence may lead to a greater understanding of the dynamical importance of turbulence to the large-scale mean flow of a turbulent fluid, we expect that, quite broadly speaking, models for MHD turbulence that are to be of some use in predicting the behavior of the mean flow must address directly the anisotropy of the turbulent Maxwell stresses, irrespective of the presence or absence of any mean field.  Many of the more recent simulations of MHD turbulence in a shearing envoronment (our main interest in this paper) are simulations (e.g. \\citet{HaGaBa:1995}) of the Velikov-Chandrasekhar-Balbus-Hawley magnetorotational instability (MRI), and we will refer to the results of these simulations for guidance. See also \\citet{Will:2002} for a comparison of simulations of the MRI with some viscoelastic models.  For our purposes, the dominant feature of simulations of the MRI is that the largest component of the turbulent stress tensors is {\\em not} the ``viscous'' $r\\theta$ component, but rather the streamwise $\\theta\\theta$ component. (Note that in simulations of the MRI, $\\theta$ is the direction of the shear, and $r$ is the cross-shear direction.) In fact, our original reason for investigating viscoelastic models was as a mechanism to collimate and possibly drive jets using this $\\theta\\theta$ component of the stress tensor.  Most importantly we wish to point out that the interpretation of the turbulent stress as an effective viscous stress can lead to qualitatively erroneous predictions. In view of the fact pointed out above that the ``viscous'' component is in fact not even the largest component of the stress tensor, a prescription for the turbulent stress tensor $\\Pi^{turb}$ of the form \\beq \\Pi^{turb}_{ij} = \\mu_{\\rm turb} \\left[\\p_i v_j + \\p_j v_i - {2 \\over 3} (\\p_k v_k) \\delta_{ij} \\right] + \\mu^{[b]}_{\\rm turb} (\\p_k v_k) \\delta_{ij}, \\eeq where $\\mu_{\\rm turb}$ and $\\mu^{[b]}_{\\rm turb}$ are effective shear and bulk viscosities respectively, is untenable, even as a gross approximation.  One possible approach to the problem is to generalize the viscosity to a tensor viscosity. This approach ignores the fact, however, that the large azimuthal component of stress in MRI-driven shear turbulence is fundamentally due to the vector advection properties of the induction equation. ",
        "conclusions": ""
    },
    "0212/astro-ph0212027_arXiv.txt": {
        "abstract": "GAIA has been approved to provide the data needed to quantify the formation and evolution of the Milky Way Galaxy, and its near neighbours. That requires study of all four key Galactic stellar populations: Bulge, Halo, Thick Disk, Thin Disk. The complex analysis methodologies required to model GAIA kinematic data are being developed, and in the interim applied to the relatively simple cases of the satellite dSph galaxies. These methodologies, illustrated here, show that we will be able to interpret the GAIA data. They also quantify what data GAIA must provide.  It is very unlikely in the present design that GAIA will be able to provide either radial velocities or worthwhile photometry for study of two of the key science goals: the Galactic Bulge and the (inner) Galactic old disk. The implication is that the radial velocity spectrometer and the medium band photometer should be optimised for study of low density fields -- suitable for their low spatial resolution -- and the broad band photometry must be optimised for inner galaxy astrophysical  studies. ",
        "introduction": "\\begin{figure}[h!] \\label{fig1} \\plotfiddle{ggilmorefig1.ps}{8cm}{0}{50}{50}{-150}{00} \\caption{Galactic population angular momentum cumulative distribution functions, showing bulge/halo (solid and dash-dot adjacent curves) and thin/thick disk (dotted and long-dashed adjacent curves) dichotomy. This indicates disparate evolutionary histories, and emphasises the critical need for GAIA to study all four Galactic components.} \\end{figure}   The GAIA mission has been approved to provide for the first time a clear picture of the formation, structure, evolution, and future of the entire Milky Way. In addition, as secondary goals, GAIA will contribute to many other branches of astrophysics, especially stellar and solar system minor body astrophysics, with a valuable contribution to cosmology and fundamental physics.   This clear scientific prioritisation must drive the design, and all compromises.  Understanding the structure and evolution of the Galaxy requires three complementary observational approaches: * Carrying out a full census of all the objects in a large, representative, part of the Galaxy; * Mapping quantitatively the spatial structure of the Galaxy; * Measuring the motions of objects in three-dimensions to determine the gravitational field and the stellar orbits.  In other words, what is required are complementary measurements of distances (astrometry), photometry to determine both extinction and intrinsic stellar properties, and the radial velocities along our line of sight.  Detailed studies of the Local Group are the key specific test of our understanding of the formation and growth of structure in the Universe.  The most significant questions here relevant for GAIA include the history of most of the stars in the Galaxy: the nature of the inner Galactic disk and the Galactic Bulge: what are their age, abundance and assembly histories? Stellar studies locally of course are of critical general significance: only locally can we determine the stellar Initial Mass Function directly.  This function directly controls the chemical and luminosity evolution of the Universe. Galactic satellite galaxies are proving the most suitable environs to quantify the nature and distribution of dark matter, and to test the small scale predictions of hierarchical galaxy formation models. The Galactic disk itself, of course, is a key test of angular momentum distributions (figure~1), chemical evolution, and merger histories, and must eventually provide the most robust information on Cold Dark Matter. Such issues as the number of local dwarf galaxies and the inner CDM profiles of the dSphs are well known demanding challenges for CDM models. Other Local group information is also important: what is the stellar Initial Mass Function, what is the distribution of chemical elements, what is the age range in the Galactic Bulge, when did the last significant disk merger happen...? All these challenge our appreciation of galaxy formation and evolution, and in turn provide the information needed to refine the models.  The partnership between local observations and {\\sl ab initio} theory is close, and developing well, and must culminate in GAIA.   There are clear similarities and distinctions between fundamental properties of the different Galactic stellar populations, such as age, metallicity, star formation history, angular momentum (figure~1) which allow study of their individual histories. This is arguably one of the greatest advantages of studies of Local Group galaxies: one is able to disentangle the many different histories which have led to a galaxy typical of those which dominate the luminosity density of the Universe. ",
        "conclusions": ""
    },
    "0212/astro-ph0212211_arXiv.txt": {
        "abstract": "Close to 70 radio pulsars have now been discovered in 23 globular clusters, with a record 22 pulsars observed in 47~Tuc alone. Accurate timing solutions, including positions in the cluster, are known for many of the pulsars. These recent observations provide a unique opportunity to re-examine theoretically the formation and evolution of recycled pulsars in globular clusters. This brief review focuses on dynamical exchange interactions between neutron stars and primordial binaries, through which neutron stars can acquire intermediate-mass binary companions. The later evolution of these intermediate-mass binaries leads naturally to the two main types of binary millisecond pulsars observed in clusters. ",
        "introduction": "The properties of globular cluster pulsars (see articles by Lorimer et al.\\ and Ransom et al.\\ in this volume) are rather surprising. While some pulsars are single, the majority are in short-period binaries. Most of the binaries have properties similar to those of the rare ``eclipsing binary pulsars'' seen in the Galactic disk population (see Nice et al.\\ 2000 for a review). These systems have extremely short orbital periods, $P_b\\sim1-10\\,$hr, circular orbits, and very low-mass companions, with $m_2\\simeq 0.03\\,M_\\odot$. The other, ``normal'' binaries have properties more similar to those of the disk population of binary millisecond pulsars, with nearly-circular orbits, periods $P_b\\sim1\\,$d (near the short-period end of the distribution for binaries in the disk) and white-dwarf (WD) companions with $m_2\\sin i\\simeq 0.2\\,M_\\odot$.  The large inferred total population of millisecond pulsars in globular clusters (several hundred in 47~Tuc alone; see Camilo et al.\\ 2000) and the very high stellar densities in many cluster cores ($\\rho_c\\sim 10^4-10^6\\,M_\\odot\\,{\\rm pc}^{-3}$) suggest that dynamical interactions must play a dominant role in the formation of these systems. However, the dynamical formation scenarios traditionally invoked for the production of recycled pulsars in globular clusters have many problems.  Scenarios based on {\\em tidal capture\\/} of low-mass main-sequence (MS) stars by neutron stars (NS), followed by accretion and recycling of the NS during a stable mass-transfer phase, run into many difficulties. Serious problems have been pointed out about the tidal capture process itself (which, because of strong nonlinearities in the regime relevant to globular clusters, is far more likely to result in a merger than in the formation of a detached binary; see, e.g., Kumar \\& Goodman 1996; McMillan et al.\\ 1990; Rasio \\& Shapiro 1991). Moreover, the basic predictions of tidal capture scenarios are at odds with many observations of binaries and pulsars in clusters (Bailyn 1995; Johnston et al.\\ 1992; Shara et al.\\ 1996). It is likely that ``tidal-capture binaries'' are either never formed, or contribute negligibly to the production of recycled pulsars. Physical {\\em collisions\\/} between NS and red giants have also been invoked as a way of producing directly NS--WD binaries with ultra-short periods (e.g., Verbunt 1987), but detailed hydrodynamic simulations show that this does not occur (Rasio \\& Shapiro 1991).  The viability of tidal capture and two-body collision scenarios has become less relevant with the realization over the last decade that globular clusters contain dynamically significant populations of {\\em primordial binaries\\/} (Hut et al.\\ 1992). Dynamical interactions involving hard primordial binaries are now thought to provide the dominant energy production mechanism that allows many globular clusters to remain in thermal equilibrium and avoid core collapse over $\\sim10^{10}\\,$yr (Gao et al.\\ 1991; McMillan \\& Hut 1994; Fregeau et al.\\ 2003).  In dense clusters, retained NS can acquire binary companions through {\\em exchange interactions\\/} with these primordial binaries. Because of its large cross section, this process dominates over any kind of two-body interaction even for low primordial binary fractions (Heggie, Hut, \\& McMillan 1996; Leonard 1989; Sigurdsson \\& Phinney 1993). In contrast to tidal capture, exchange interactions with hard primordial binaries (with semimajor axes $a\\sim0.1-1\\,$AU) can form naturally the wide binary millisecond pulsars seen in some low-density globular clusters (such as PSR B1310$+$18, with $P_b=256\\,$d, in M53, which has the lowest central density, $\\rho_c\\sim10^3\\,M_\\odot\\,{\\rm pc}^{-3}$, of any globular cluster with observed radio pulsars). When the newly acquired MS companion, of mass $\\la1\\,M_\\odot$, evolves up the giant branch, the orbit circularizes and a period of {\\em stable\\/} mass transfer begins, during which the NS is recycled (see, e.g., Rappaport et al.\\ 1995). The resulting NS--WD binaries have orbital periods in the range $P_b\\sim1-10^3\\,$d. However, this scenario cannot explain the formation of recycled pulsars in binaries with periods shorter than $\\sim1\\,$d. To obtain such short periods, the initial primordial binary must be extremely hard, with $a \\la 0.01\\,$AU, but then the recoil velocity of the system following the exchange interaction would almost certainly exceed the escape speed from the shallow cluster potential (e.g., $v_e\\simeq 60\\,{\\rm km}\\,{\\rm s}^{-1}$ for 47~Tuc).  One can get around this problem by considering more carefully the stability of mass transfer in NS--MS binaries formed through exchange interactions. While all MS stars in the cluster {\\em today\\/} have masses $\\la1\\,M_\\odot$, the rate of exchange interactions should have peaked at a time when significantly more massive MS stars were still present. Indeed, the NS and the most massive primordial binaries will undergo mass segregation and concentrate in the cluster core on a time scale comparable to the initial half-mass relaxation time $t_{\\rm rh}$. For typical dense globular clusters, we expect $t_{\\rm rh}\\sim 10^9\\,$yr, which is comparable to the MS lifetime of a $\\sim 2-3\\,M_\\odot$ star. If the majority of NS in the cluster core acquire MS companions in the range of $\\sim1-3\\,M_\\odot$, a very different evolution ensues. Indeed, in this case, when the MS star evolves and fills its Roche lobe, the mass transfer for many systems (depending on the mass ratio and evolutionary state of the donor star) is {\\em dynamically unstable\\/} and leads to a common-envelope (CE) phase (see, e.g., Taam \\& Sandquist 2000). The emerging binary will have a low-mass WD in a short-period, circular orbit around the NS.  This simple idea has been explored quantitatively using Monte Carlo simulations of globular cluster dynamics (Rasio, Pfahl, \\& Rappaport 2000; Rappaport et al.\\ 2001) and a brief summary of this work is presented in Sec.~2. A similar scenario, but starting from tidal capture binaries and applied to X-ray sources in globular clusters, was discussed by Bailyn \\& Grindlay (1987). The possibility of forming intermediate-mass binaries through exchange interactions was mentioned by Davies \\& Hansen (1998), who pointed out that NS retention in globular clusters may also require that the NS be born in massive binaries. Among eclipsing pulsars in the disk, at least one system (PSR J2050$-$0827) is likely to have had an intermediate-mass binary progenitor, given its very low transverse velocity (Stappers et al.\\ 1998). ",
        "conclusions": ""
    },
    "0212/astro-ph0212161_arXiv.txt": {
        "abstract": "AA\\,Dor\\index{AA Dor} (LB\\,3459\\index{LB 3459}) is an eclipsing, close binary ($P = 0.26\\,\\mathrm{d}$) consisting of a sdOB primary star and an unseen secondary with an extraordinary small mass. The secondary is possibly a former planet which may have survived a common-envelope phase and has even gained mass.  In order to investigate on an existing discrepancy between the components' mass derived from NLTE spectral analysis and subsequent comparison to evolutionary tracks and masses derived from radial-velocity and the eclipse curves, we performed phase-resolved high-resolution and high-SN spectroscopy with the UVES attached to the ESO VLT. From the obtained spectra, we have determined AA\\,Dor\\index{AA Dor}'s orbital parameters ($P = 22\\,600.702 \\pm 0.005 \\mathrm{sec}$, $A = 39.19 \\pm 0.05\\mathrm{km/sec}$, $T_0 = 2451917.152690$) and the rotational velocity ($v_{\\mathrm rot} = 47 \\pm 5\\mathrm{km/sec}$) of the primary. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/astro-ph0212482_arXiv.txt": {
        "abstract": " ",
        "introduction": "The estimation of variations in CNO abundances within globular clusters yields information on the evolutionary history of the clusters, and possibly also on the physical processes involved in their formation.  This paper describes some preliminary results from a new photometric method to estimate these variations. The method can be used for any globular cluster, but M22 was chosen for a first investigation because it is known (Richter, Hilker, \\& Richtler 1999) to exhibit large internal variations in CN absorption strength, and these variations, unlike in most clusters, are believed (Anthony-Twarog, Twarog, \\& Craig 1995) to be {\\it positively} correlated with CH variations. ",
        "conclusions": ""
    },
    "0212/astro-ph0212357_arXiv.txt": {
        "abstract": "The halo occupation distribution (HOD) describes the bias between galaxies and dark matter by specifying (a) the probability $\\PNM$ that a halo of virial mass $M$ contains $N$ galaxies of a particular class and (b) the relative spatial and velocity distributions of galaxies and dark matter within halos.  We calculate and compare the HODs predicted by a smoothed particle hydrodynamics (SPH) simulation of a $\\Lambda$CDM cosmological model (cold dark matter with a cosmological constant) and by a semi-analytic galaxy formation model applied to the same cosmology. Although the two methods predict different galaxy mass functions, their HOD predictions for samples of the same space density agree remarkably well. In a sample defined by a baryonic mass threshold, the mean occupation function $\\N_M$ exhibits a sharp cutoff at low halo masses, a slowly rising plateau in which $\\N$ climbs from one to two over nearly a decade in halo mass, and a more steeply rising, high occupancy regime at high halo mass. In the low occupancy regime, the factorial moments $\\NN$ and $\\NNN$ are well below the values $\\N^2$ and $\\N^3$ expected for Poisson statistics, with important consequences for the small scale behavior of the 2- and 3-point correlation functions. The HOD depends strongly on galaxy age, with high mass halos populated mainly by old galaxies and low mass halos by young galaxies. The distribution of galaxies within SPH halos supports the assumptions usually made in semi-analytic calculations: the most massive galaxy lies close to the halo center and moves near the halo's mean velocity, while the remaining, satellite galaxies have the same radial profile and velocity dispersion as the dark matter. The mean occupation at fixed halo mass in the SPH simulation is independent of the halo's larger scale environment, supporting both the merger tree approach of the semi-analytic method and the claim that the HOD provides a complete statistical characterization of galaxy bias. We discuss the connections between the predicted HODs and the galaxy formation physics incorporated in the SPH and semi-analytic approaches. These predictions offer useful guidance to theoretical models of galaxy clustering, and they will be tested empirically by ongoing analyses of galaxy redshift surveys. By applying the HODs to a large volume N-body simulation, we show that both methods predict slight departures from a power-law galaxy correlation function, similar to features detected in recent observational analyses. ",
        "introduction": "\\label{intro}  A complete theory of galaxy formation should predict the distributions of galaxy luminosities, colors, sizes, and morphologies, the correlations among these properties, and the relation between the spatial clustering of any given class of galaxies and that of the underlying dark matter distribution. This last class of predictions, the ``bias'' of galaxies as a function of their observable properties, is becoming an increasingly important test of theoretical models thanks to the new generation of large galaxy redshift surveys, in particular the 2dF Galaxy Redshift Survey (2dFGRS; \\citealt{colless01}) and the Sloan Digital Sky Survey (SDSS; \\citealt{york00}). The ``halo occupation distribution'' (HOD) formalism is an especially powerful framework for carrying out such tests; it characterizes the bias of a class of galaxies by the probability $\\PNM$ that a halo of virial mass $M$ contains $N$ such galaxies and additional prescriptions that specify the relative distributions of galaxies and dark matter within halos. If the HOD at fixed halo mass is statistically independent of the halo's large scale environment, as theoretical models predict (\\citealt{bond91,white96,lemson99}; this paper), then this description of galaxy bias is essentially complete: given the HOD and the halo population predicted by a particular cosmological model, one can calculate any galaxy clustering statistic, on scales from the linear regime to the deeply non-linear regime. Empirical determinations of the HOD for different galaxy types would therefore summarize everything that observed galaxy clustering has to say about the physics of galaxy formation, in a form that can be readily compared to theoretical predictions.  This paper examines the HODs predicted by the two leading theoretical methods for studying galaxy formation and bias in a cosmological context: semi-analytic models (e.g., \\citealt{white91,kauffmann93,cole94,avila98,somerville99}) and hydrodynamic numerical simulations (e.g., \\citealt{cen92,katz92,evrard94,pearce99,white01,yoshikawa01}). We apply both methods to a $\\Lambda$CDM cosmological model (inflationary cold dark matter with a cosmological constant), adopting the same cosmological parameters in each case.  We focus on galaxy samples defined by thresholds in baryon mass and stellar population age, which are roughly analogous to observational samples defined by cuts in luminosity and color. Except for using the same cosmological model and thus the same present-day halo population, we do not make any efforts to ``tune'' the semi-analytic calculation to match the hydrodynamic simulation; we apply each method in its ``standard'' form. We present results for the various features of the HOD --- the mean occupation $\\N$ as a function of halo mass, moments of the distribution $\\PNN$, and the spatial and velocity distributions of galaxies within halos --- and we interpret these features in terms of the physical processes represented in the theoretical models. In the short term, our results should provide useful input to theoretical models of galaxy clustering by allowing the predictions of these galaxy formation models to be ``bootstrapped'' into analytic calculations or larger volume N-body simulations, and they offer guidance to efforts to infer parameters of the HOD from observational data. In the slightly longer term, these predictions will be tested by empirical determinations of the HOD, and any discrepancies with observations may point the way to necessary revisions of the galaxy formation model or the underlying cosmological model.  Several aspects of the HOD predicted by semi-analytic models have been investigated in the pioneering papers of \\cite{kauffmann97}, \\cite{governato98}, \\cite{kauffmann99}, and \\cite{benson00}, which used semi-analytic methods to assign galaxy populations to the halos of N-body simulations.  \\cite{seljak00} and \\cite{sheth01a} measured $\\PNM$ from the \\cite{kauffmann99} models and incorporated them into analytic predictions of galaxy clustering and bias (see also \\citealt{scoccimarro01,sheth01,scranton02}). HOD predictions of hydrodynamic simulations have been presented by \\cite{white01} at redshifts $z=3$ and $z=1$ and by \\cite{yoshikawa01} at $z=3$, 2, and 0. Relative to this earlier work, our examination of the HOD in this paper is more comprehensive, and our side-by-side comparison of numerical and semi-analytic results for the same cosmological model allows us to evaluate the robustness and limitations of the predictions and to better understand the physics that gives rise to them.  Other studies of the halo occupation distribution have focused on the connections between the HOD and statistical measures of galaxy clustering. Many of these studies have utilized the power of the HOD formalism and the related ``halo model'' of dark matter clustering as a tool for analytic calculations \\citep{ma00,seljak00,benson01,scoccimarro01,sheth01,white01b,cooray02b}, drawing on methods developed over the course of several decades \\citep{neyman52,peebles74,mcclelland77,scherrer91,mo96,sheth01b}. Berlind \\& Weinberg (\\citeyear{berlind02}, hereafter BW) computed the impact of HOD bias on many of the most widely used galaxy clustering statistics by applying parameterized HOD models to an N-body simulation of the $\\Lambda$CDM scenario.  They concluded that different statistics constrain the HOD in complementary ways, making it possible to determine the HOD empirically from observed galaxy clustering, at least for a known cosmological model.  Important steps toward observational determination of the HOD of bright, optically selected galaxies, drawing mainly on the 2- and 3-point correlation functions, the group multiplicity function, and galaxy-galaxy lensing, have been taken by \\cite{jing98}, \\cite{peacock00}, \\cite{scoccimarro01}, \\cite{guzik02}, \\cite{marinoni02}, \\cite{yang03}, \\cite{vandenbosch02}, \\cite{zehavi03}, and \\cite{magliocchetti03}.  \\cite{kochanek03}, \\cite{jing02}, and \\cite{cooray02} have investigated HOD constraints for galaxies selected in the near- or far-infrared, and \\cite{wechsler01}, \\cite{bullock02}, and \\cite{moustakas02} have applied similar methods to high-redshift galaxies. An ambitious but, we think, realizable goal is to use the high-precision measurements afforded by the 2dFGRS and the SDSS to break the ``degeneracies'' between cosmology and bias, obtaining tight, simultaneous constraints on the mass function and clustering of dark halos and the HODs of many different galaxy classes (see discussions by BW, \\citealt{zheng02}, and \\citealt{weinberg02b}). We hope that the results presented here will provide inspiration to such efforts, by illustrating how measurements of the HOD can test basic ideas about the physics of galaxy formation.  Our approach to the HOD is essentially the one described by BW, which was in turn inspired largely by the discussion of Benson et al. (2000ab). The clustering of galaxies predicted by the semi-analytic model and hydrodynamic simulation investigated here, as quantified by more traditional statistics, has been presented in separate papers \\citep{benson00,weinberg02a}. We briefly describe the two calculations and the selection of galaxy populations in \\S~\\ref{models}. In \\S~\\ref{pnm} we compare the predictions of $\\PNM$ for several different galaxy classes, and we demonstrate that $\\PNM$ predicted by the hydrodynamic simulation is independent of large scale environment, to within the statistical limitations of our measurements. Semi-analytic models do not predict the distribution of galaxies within halos, but clustering calculations based on these models usually assume that each halo contains one central galaxy moving at the halo's center-of-mass velocity and that other ``satellite'' galaxies trace the halo's dark matter distribution. In \\S~\\ref{profile}, we show that these assumptions hold to a good approximation in the hydrodynamic simulation, which predicts the galaxy positions and velocities directly. In \\S~\\ref{central} we compare the properties of central and satellite galaxies in the two methods. The precision in predictions of the galaxy correlation function by these two models has been limited by the finite size of the simulation volumes in which they were implemented. In \\S~\\ref{xiJen} we combine the HOD results with a new, large-volume N-body simulation to make improved predictions for the galaxy correlation function, focusing on predicted departures from a power-law form. In \\S~\\ref{conclusions} we summarize our results and discuss what they tell us about the physical factors that shape the HOD and thereby determine the bias between galaxies and dark matter. ",
        "conclusions": "\\label{conclusions}  The SPH and SA predictions of the galaxy HOD agree remarkably well. The qualitative features of these predictions can, for the most part, be readily interpreted in terms of galaxy formation physics.  The mean occupation function for galaxies above a mass threshold has a cutoff, a slowly rising plateau, and a steeper ``high occupancy'' regime.  In the SPH simulation, the cutoff is determined by the universal baryon fraction. In the SA model, the cutoff occurs at a substantially higher mass than the universal baryon fraction requires, corresponding to only $\\sim$25\\% of the gas in low mass halos remaining in galaxies.  This is due to the feedback mechanism, which in the SA model is tuned to produce the observed galaxy luminosity function.  The cutoff is nearly as sharp in the SA model as in the SPH simulation because feedback is tightly regulated by halo mass, with no scatter that allows some low mass halos to retain substantially more than $\\sim$25\\% of their baryons.  In the low occupancy regime, $\\N$ rises slowly with $M$ because additional mass tends to go into making more massive single galaxies instead of multiple low mass galaxies.  The logarithmic slope of $\\N_M$ steepens for $\\N \\gtrsim 2$, presumably representing a transition to a regime in which the galaxies in a halo were formed mostly in the lower mass progenitor halos that merged to create it.  Even in this regime, the slope is below unity, since the overall efficiency of galaxy formation (fraction of mass in galaxies) is lower in high mass halos, probably a consequence of longer cooling times and less efficient filamentary ``channeling'' of gas into galaxies. For higher baryonic mass thresholds, the mean occupation function shifts horizontally along the $\\log M$ axis, so this physical explanation of its form is not sensitive to the particular threshold adopted.  Since the considerations that account for the form of $\\N_M$ are quite general, it is perhaps not surprising that the SPH and SA calculations agree.  More impressive is the agreement on fluctuations about the mean occupation.  Both calculations predict pair counts well below the expectation for Poisson fluctuations at low mass, close to those of the minimal fluctuation, nearest-integer model.  As many authors have emphasized (\\citealt{benson00,seljak00,peacock00,scoccimarro01}; BW), these sub-Poisson fluctuations are essential in understanding the power-law form of the observed galaxy correlation function.  Triple counts are also substantially sub-Poisson, though not minimal, and the two calculations again agree. The explanation of sub-Poisson fluctuations is not trivial, but it clearly has to do with statistics of halo merger histories and the timescales of galaxy mergers following halo mergers.  In particular, near equal-mass halo mergers are rare enough that one must go substantially above the cutoff mass $\\Mmin$ (by a factor $\\sim10$) before one is likely to get two galaxies above the mass threshold in a single halo, instead of one, more massive, galaxy. The probability of empty halos follows the nearest-integer prediction, $P(0|M)={\\rm max}(1-\\N_M,0)$, almost exactly in both models. Because the cutoff of $\\N_M$ is somewhat sharper in the SPH calculation, the transition from occupied halos to empty halos (at the specified baryonic mass threshold) is also somewhat sharper, occurring over about a factor of two in halo mass.  For masses above the cutoff, the probability that a halo is empty is much lower than the Poisson expectation $\\exp(-\\N_M)$.  In this regime, strongly sub-Poisson statistics indicate the regularity of the formation process in single-galaxy halos, which produces a tight correlation between galaxy mass and halo mass.  The mean occupation is very different for galaxies in different age quartiles, at both the low and high mass ends of $\\N_M$. Young galaxies are rare in high mass halos and vice versa. This difference reflects a tendency for star formation to begin early in regions that will become part of high mass halos and for star formation to shut off when a galaxy falls into a halo of larger virial mass, or to slow down when the halo size or virial temperature becomes too large. The qualitative agreement between the SPH and SA calculations is again not surprising, but the degree of quantitative agreement is impressive.  The SPH simulation also shows that the mean occupation function is independent of the large scale environment of the halo, extending to the galaxy regime the result that \\citet{lemson99} found for the merger histories of dark matter halos.  This is a fundamental and encouraging result, supporting the assumption made in merger-tree based SA methods, and supporting the claim that the HOD provides a statistically complete description of galaxy bias.  A detailed analysis of the SPH simulation also supports simple assumptions about the galaxy distribution within halos.  Most halos have a galaxy near the center of mass moving at close to the center-of-mass velocity. The central galaxy is almost always the most massive galaxy in the halo, and it is usually among the oldest galaxies in the halo. The remaining, satellite galaxies have the same spatial and velocity distribution as the dark matter, except for a small ($\\alpha_v \\sim 0.8-0.95$) velocity bias in intermediate mass halos. These assumptions are thus reasonable to use when making galaxy clustering predictions, and they are the ones usually incorporated in studies that combine N-body and semi-analytic methods (\\citealt{kauffmann99}; \\citealt{benson00}; \\citealt{somerville01}). In the long term, these predictions can be tested empirically using clustering and galaxy-galaxy lensing data of the sort provided by 2dFGRS and SDSS.  As we have emphasized in separate papers on the SA model and SPH simulation (\\citealt{benson00}; \\citealt{weinberg02a}), the complex relation between galaxies and mass is crucial in reproducing the observed nearly power-law form of the galaxy correlation function, since the dark matter $\\xi(r)$ is not a power law in CDM-type models.  Since the emergence of the power law is to some extent a coincidence, the HOD framework generically predicts that there should be departures from a power-law $\\xi(r)$ once it is measured with sufficient precision, especially if one considers different subclasses of galaxies.  We have tried to predict these departures by bootstrapping our results into a larger volume N-body simulation. The small differences between the SA and SPH $\\PNM$ translate into noticeable differences in the predicted $\\xi(r)$, and the accuracy of the SPH prediction is still limited to some extent by the small number of high mass halos in the SPH volume.  However, improved statistics show that both models predict an inflection in $\\xi(r)$ at $r\\sim 1\\hmpc$, with the departures from a power law being stronger for the SPH model than for the SA model and stronger for more massive galaxies in both cases.  Recent observational analyses provide evidence for just such a feature in the correlation function of luminous galaxies \\citep{zehavi03}.  The most surprising aspect of our results is that the SPH and SA models give such similar predictions for $\\PNM$ despite having very different galaxy mass functions. The different mass functions reflect the different treatments of cooling in the SA and SPH calculations (including the geometric idealizations in the former and the numerical resolution limitations in the latter), and the greater importance of stellar feedback in the SA calculation.  The similarity of the predicted halo occupations implies that the HOD is determined by physics that is robust to these cooling and feedback differences.  Obviously it is crucial that we select galaxy samples based on number density, using only the rank order of the baryonic masses rather than their absolute values; for a fixed mass threshold, the numbers of galaxies would be different in the two calculations, and the HODs could not possibly match.  Similarity of HODs suggests that the rank order of galaxy masses and the rank order of stellar population ages are related in fairly simple ways to the merger histories of the dark matter halos, which should be statistically similar in the two methods. For example, the baryonic mass of a galaxy could be monotonically related (by non-linear functions that are different in the two calculations) to the mass of its parent halo at the last time this galaxy was ``central.'' The age could be determined (at least in a rank-order sense) by the formation time of this halo and the time at which its central object's star formation is suppressed by falling into a larger halo and becoming a satellite.  Mergers within halos also affect $\\PNM$, but this process should be similar in the two calculations, since the assumptions about mergers in the SA model reproduce the results of numerical calculations \\citep{benson02}.  We can test this idea --- that the form of the HOD is driven largely by dark matter dynamics --- by re-analyzing the SA galaxy population without using the predicted baryonic properties. For every SA galaxy, we identify its parent halo at the last time this galaxy was central, and we denote the halo mass at that time by $\\Mmax$ and the time itself by $\\tmax$. We then create galaxy samples with the same space densities as before but use thresholds in $\\Mmax$ rather than in baryonic mass. Dotted curves in Figures~\\ref{fig:Nn} and~\\ref{fig:NNn} show the first and second moments of $N(M)$, respectively, for these samples defined by $\\Mmax$ thresholds. It is interesting both that this simple mapping works as well as it does and that it does not work perfectly. The qualitative similarity of the mean occupations supports the idea that $\\PNM$ is largely governed by robust physics connected to halo merger histories.  However, obvious differences, in particular the steeper slope of $\\N_M$ at the high mass end, indicate that additional physical processes must also play a significant role. It is encouraging that the SPH and SA models agree with each other better than either model agrees with the ``rank by $\\Mmax$'' calculation. The fact that they predict fewer galaxies in high mass halos suggests that galaxy accretion rates begin to drop in overdense environments even before their parent halos fall into larger halos, allowing some of these galaxies to drop below the baryonic mass threshold and be replaced (in a sample of fixed space density) by galaxies forming in lower mass halos. Recent studies provide observational support for this effect, showing decreased star formation rates in galaxies that are well outside the virial radii of rich clusters (e.g., \\citealt{lewis02,gomez03}).  We can map the stellar age of a galaxy onto the halo merger tree in a similar fashion.  We make the simple assumption that the star formation in each galaxy starts at a time $\\tstart$ when its parent halo first exceeds a mass $10^{11}\\Msun$, a factor $\\sim 4$ below the cutoff for $\\ng=0.02\\hvol$.  We assume that the star formation ends at a time $\\tstop$ when the galaxy becomes a satellite ($\\tstop=\\tmax$) or when the cooling time in its parent halo becomes too long ($\\tstop=t_{\\rm hot}$) --- whichever happens first.  We identify $t_{\\rm hot}$ as the time when the parent halo mass exceeds $10^{13}\\Msun$.  Finally, we assume a constant star formation rate between $\\tstart$ and $\\tstop$, making the mean stellar age of the galaxy $t_* = t_0 - (\\tstop-\\tstart)/2$.  Figure~\\ref{fig:tstar} compares the mean occupation functions for SA galaxies split into quartiles according to $t_*$ to those of quartiles defined by true stellar age (repeated from Figure~\\ref{fig:Nage}).  Our highly simplified approximation yields the right qualitative dependence of $\\N_M$ on galaxy age, but it differs substantially in the quantitative details.  The $\\Mmax$ and $t_*$ models still make indirect use of the baryonic physics in the SA calculation, since baryonic masses affect which galaxies remain as distinct entities in larger halos and which are destroyed by dynamical friction.  Results from a more concerted effort to model the HOD entirely with dark matter dynamics will be reported elsewhere (Berlind, Bullock, Kravtsov, Wechsler, and Zentner in preparation).  \\begin{figure} \\plotone{f23.eps} \\caption{Mean occupation $\\N_M$ of SA galaxies, for quartiles of mean stellar galaxy age (solid curves) and merger tree ``age'' parameter $t_*$ (dashed curves, defined in \\S~\\ref{conclusions}). } \\label{fig:tstar} \\end{figure}  Finite resolution of the SPH simulation has limited our comparison of predictions to relatively massive galaxies, and to age rather than morphology as a distinguishing characteristic of sub-classes. Insensitivity of predictions to calculational details is likely to change when one gets to low mass galaxies or to characteristics that are determined by interactions with environment.  For example, we might expect the HOD of low mass galaxies to be quite different if their luminosity is controlled by photoionization or by supernova feedback.  Likewise, models that ascribe bulge-to-disk ratios mainly to merger histories, to secular evolution of disks, to weak perturbations in group/cluster environments, or to interaction with the IGM could yield quite different predictions for the $\\PNM$ of different morphological types. The process of inferring halo occupations from observed galaxy clustering has already begun, and it should accelerate over the next few years with improved measurements from the 2dFGRS and the SDSS. Empirical HOD determinations for galaxies with $L\\ga L_*$, classified by luminosity and color or spectral type, will test the basic predictions of the current scenario of galaxy formation, as presented here.  They will also sharpen the constraints on cosmological models by removing bias as a degree of freedom in matching observed clustering. Empirical determinations of HODs for faint galaxies and for morphological subtypes should yield insight into the physics that governs the low end of the luminosity function and the origin of galaxy morphology."
    },
    "0212/astro-ph0212577_arXiv.txt": {
        "abstract": "s{ We describe a newly developed cosmological hydrodynamics code based on the weighted essentially non-oscillatory (WENO) schemes for hyperbolic conservation laws. High order finite difference WENO schemes are designed for problems with piecewise smooth solutions containing discontinuities, and have been successful in applications for problems involving both shocks and complicated smooth solution structures. We couple hydrodynamics based on the WENO scheme with standard Poisson solver - particle-mesh (PM) algorithm for evolving the self-gravitating system.  A third order total variation diminishing (TVD) Runge-Kutta scheme has been used for time-integration of the system. We brief the implementation of numerical technique. The cosmological applications in simulating intergalactic medium and Ly$\\alpha$ forest in the CDM scenario are also presented. } ",
        "introduction": "Due to the highly non-linearity of gravitational clustering in the universe, one significant feature emerging in cosmological hydrodynamic flow is extremely supersonic motion around the density peaks developed by gravitational instability, which leads to strong shock discontinuities within complex smooth structures. It poses more challenge than the typical hydrodynamic simulation without self-gravity. To address this issue, two algorithms have been implemented for high resolution shock capturing in cosmology: total-variation diminishing (TVD) scheme$^{[1]}$ and piecewise parabolic method (PPM)$^{[2]}$.  An alternative hydrodynamic solver to discretize the convection terms in the Euler equations is the fifth order finite difference WENO (weighted essentially non-oscillatory) method$^{[3]}$, with a low storage nonlinearly stable third order Runge-Kutta time discretization$^{[4]}$. The WENO schemes, first constructed in third order finite volume version by Liu et al.$^{[5]}$, are based on the essentially non-oscillatory scheme (ENO) developed by Harten et al.$^{[6]}$ in the form of finite volume scheme for hyperbolic conservative laws. The ENO scheme generalizes the total variation diminishing (TVD) scheme of Harten. The TVD schemes typically degenerate to first-order accuracy at locations with smooth extrema while the ENO and WENO schemes maintain high order accuracy there even in multi-dimensions. The WENO method is very robust and stable for solutions containing strong shocks and complex solution structures$^{[3],[5]}$. It uses the idea of adaptive stencils in the reconstruction procedure based on the local smoothness of the numerical solution to automatically achieve high order accuracy and non-oscillatory property near discontinuities. This is achieved by using a convex combination of a few candidate stencils, each being assigned a nonlinear weight which depends on the local smoothness of the numerical solution based on that stencil.  WENO schemes can simultaneously provide a high order resolution for the smooth part of the solution, and a sharp, monotone shock or contact discontinuity transition.  In the context of cosmological applications, we have developed a hybrid N-body/hydrodynamical code that incorporates a Lagrangian particle-mesh algorithm to evolve the collision-less matter with the fifth order WENO scheme to solve the equation of gas dynamics. This paper describes this code briefly. ",
        "conclusions": ""
    },
    "0212/astro-ph0212094_arXiv.txt": {
        "abstract": "{We examined the Fourier power spectrum characteristics of cirrus structures in 13 sky fields with faint to bright cirrus emission observed with ISOPHOT in the 90--200$\\mu$m wavelength range in order to study variations of the spectral index $\\alpha$. We found that $\\alpha$ varies from field to field with --5.3\\,$\\le$\\,$\\alpha$\\,$\\le$\\,--2.1. It depends on the absolute surface brightness and on the hydrogen column density. We also found different spectral indices for the same sky region at different wavelengths. Longer wavelength measurements show steeper power spectra. This can be explained by the presence of dust at various temperatures, in particular of a cold extended component. For the faintest areas of the far-infrared sky we derive a wavelength independent spectral index of $\\alpha$\\,=\\,--2.3$\\pm$0.6  for the cirrus power spectrum. The application of the correct spectral index is a precondition for the proper disentanglement of the cirrus foreground component of the Cosmic Far-Infrared Background and its fluctuations. ",
        "introduction": "} \\begin{figure*} \\hbox{\\hskip 1.1cm \\epsfxsize=4.8cm \\epsffile{ms2481f1a.ps} \\hskip 1.2cm \\epsfxsize=4.8cm \\epsffile{ms2481f1c.ps}} \\vskip 0.8cm \\hbox{\\hskip 1.1cm \\epsfxsize=4.8cm \\epsffile{ms2481f1b.ps} \\hskip 1.2cm \\epsfxsize=4.8cm \\epsffile{ms2481f1d.ps}} \\vskip -5.8cm \\hfill \\parbox[b]{48mm}{ \\caption[]{100 and 200\\,$\\mu$m range images of two sample fields in the satellite coordinate system \\citep{Laureijs2001}. The central sky positions of the images are listed in Table~1. The orientation of the celestial coordinate system is indicated. ({\\bf a}) Cham--II, 200\\,$\\mu$m ({\\bf b}) Cham--II, 100\\,$\\mu$m ({\\bf c}) NPC\\,1, 200\\,$\\mu$m ({\\bf d}) NPC\\,1, 90\\,$\\mu$m. While the FIR emission in NPC\\,1 arises from the same cirrus structure at both wavelengths, there seems to be an additional component of cold condensations in Cham\\,II, as discussed in Sect.~\\ref{sect:disc_alphavar}.  } \\label{fig:images}} \\vskip 0.5cm \\end{figure*} In the far-infrared the most significant component superimposed on the Extragalactic Background  is the cirrus emission of the Galaxy originating from irregularly shaped interstellar clouds \\citep{Low_84}. The correct determination of the Cosmic Far-InfraRed Background (CFIRB) requires a precise removal of this foreground component {\\citep{Guiderdoni}}. The CFIRB is characterized by two parameters: the absolute level and the amplitude of its fluctuations. In recent years observations made by ISOPHOT \\citep{Lemke96}, the photometer on board the ISO satellite \\citep{Kessler96} made the determination of the CFIRB fluctuations possible.  CFIRB fluctuations are expected to show the following properties {\\citep[see e.g.][]{Hauser+Dwek}}: (1) the spatial distribution corresponds to a flat power spectrum, (2) they are isotropic and (3) they have a positive constant value. In contrast, the fluctuations due to galactic cirrus emission scale up with absolute brightness and can be described by a multifractal structure, yielding a steep Fourier power spectrum. It is usually described by one main parameter, the spectral index $\\alpha$: \\begin{equation} \\rm P = P_0 \\times (f/f_{0})^\\alpha \\label{eq:powerlaw} \\end{equation} where P is the Fourier power, f is the absolute value of the two-dimensional spatial frequency and P$_0$ is the Fourier power at the spatial frequency f$_0$ \\citep{Gautier}. Since cirrus fluctuations are different from CFIRB fluctuations in at least two features (spatial frequency-- and brightness--dependence), it is possible to separate them from each other in two independent ways:  (1) The decomposition through the brightness dependence at a fixed spatial frequency was performed by \\citet[][hereafter Paper I]{Kiss2001}, and resulted in a new determination of the CFIRB fluctuation amplitudes at 90 and 170\\,$\\mu$m.  (2) The decomposition using the Fourier power spectrum was introduced by \\citet{Guiderdoni}. For faint fields, where CFIRB fluctuations are expected to be dominant in the highest spatial frequency range, the knowledge of the cirrus spectral index is mandatory for decomposition. \\citet{Gautier} determined a spectral index of the galactic cirrus emission $\\alpha$\\,$\\approx$\\,--3 from the analysis of 100\\,$\\mu$m IRAS scans. For regions of faint total FIR brightness emission \\citet{Lagache2000} and \\citet{Matsuhara} derived CFIRB fluctuation amplitudes applying this spectral index of $\\alpha$\\,$\\approx$\\,--3 for the cirrus decomposition.  \\citet{Herbstmeier} analysed a few fields observed by ISOPHOT in the 90--180\\,$\\mu$m wavelength range. They found a variation of spectral indices in the range of --0.5\\,$\\ge$\\,$\\alpha$\\,$\\ge$\\,--3.6, depending on wavelength, brightness and sky position. However, this sample was too small (four fields, two of them measured only at one wavelength) for general conclusions, but their results clearly indicated that the $\\alpha$\\,=\\,--3 cirrus spectral index may not be universal.  Despite the strong impact of the cirrus power spectrum on the correct determination of the characteristics of the CFIRB, cirrus power spectrum properties were not yet analysed in detail for $\\lambda$\\,$>$\\,100$\\mu$m and at a resolution achievable by ISOPHOT. In this paper we analyse 13 sky regions measured by ISOPHOT in the 90--200\\,$\\mu$m range and investigate the characteristics of their Fourier power spectra. ",
        "conclusions": "} We examined the Fourier power spectrum characteristics of 13 sky fields with faint to bright cirrus emission in the 90--200$\\mu$m wavelength range and derived the spectral index of the power spectrum, $\\alpha$. We found that $\\alpha$ varies from field to field. It has a clear dependence on the absolute surface brightness and on the hydrogen column density of the field. For longer wavelengths ($\\lambda$\\,$>$\\,100\\,$\\mu$m) the scatter in $\\alpha$ is larger, especially for higher hydrogen column densities. This is because for N(HI)\\,$\\ge$\\,10$^{21}$\\,cm$^{-2}$ the atomic-to-molecular phase transition leads to various $\\alpha$ values especially in the 170--200\\,$\\mu$m range, where the diverse molecular content is easier to observe. We also found a significant difference in spectral indices for the same sky region at different wavelengths. At longer wavelengths ($\\lambda\\,>$\\,100\\,$\\mu$m) the power spectra are steeper. This can be explained by the coexistence of dust components with various temperatures within the same field and cold extended emission features which significantly affect the power spectrum. A spectral index of the cirrus power spectrum with a value of $\\alpha$\\,=\\,--2.3$\\pm$0.6 -- independent of wavelength -- could be derived for the faintest areas of the far-infrared sky. The precise determination of the spectral index for each field will enable to correctly disentangle the galactic foreground components of the CFIRB.  \\rm"
    },
    "0212/astro-ph0212049_arXiv.txt": {
        "abstract": "\\objectname[]{NGC~2915} is a blue compact dwarf galaxy embedded in an extended, low surface brightness \\ion{H}{1} disk exhibiting a two-armed spiral structure and a central bar-like component. Commonly accepted mechanisms are unable to explain the existence of these patterns and Bureau et al.\\ proposed disk dark matter (scaling with the \\ion{H}{1} distribution) or a rotating triaxial dark halo as alternative solutions. In an attempt to explore these mechanisms, hydrodynamical simulations were run for each case and compared to observations using customized column density and kinematic constraints. The spiral structure can be accounted for both by an unseen bar or triaxial halo, the former fitting the observations slightly better. However, the large bar mass or halo pattern frequency required make it unlikely that the spiral wave is driven by an external perturber. In particular, the spin parameter $\\lambda$ is much higher than predicted by current cold dark matter (CDM) structure formation scenarios. The massive disk models show that when the observed gas surface density is scaled up by a factor about $10$, the disk develops a spiral structure resembling closely the observed one, in perturbed density as well as perturbed velocity. This is consistent with more limited studies in other galaxies and suggests that the disk of \\objectname[]{NGC~2915} contains much more mass than is visible, tightly linked to the neutral hydrogen. A classic (quasi-)spherical halo is nevertheless still required, as increasing the disk mass further to fit the circular velocity curve would make the disk violently unstable. Scaling the observed surface density profile by an order of magnitude brings the disk and halo masses to comparable values within the disk radius. The surface density remains under Kennicutt's star formation threshold for a gaseous disk and no stars are expected to form, as required by observations. ",
        "introduction": "}  \\objectname[]{NGC~2915} is a blue compact dwarf (BCD) galaxy at a distance $D=5.3\\pm1.6$~Mpc \\citep*{mmc94}. In the optical, it displays a high surface brightness blue core and a red diffuse population. The core is the locus of high mass star formation and possesses a high excitation, low metallicity \\ion{H}{2} region spectrum; the red population has an exponential profile with a low extrapolated central surface brightness \\citep{mmc94}. The radio properties of \\objectname[]{NGC~2915} are rather extreme, with an \\ion{H}{1} disk extending to 22 radial scalelengths in the $B$-band \\citep[5 Holmberg radii;][]{mcbf96}. This disk displays a short central bar overlapping the optical emission and a well-developed outer two-armed spiral extending to its edge (see Fig.~\\ref{fig:iv_images}). The signature of an oval distortion is also present in the velocity field. Standard modeling of the derived circular velocity curve yields a blue mass-to-light ratio ${\\cal M}_T/L_B\\ga65$ at the last measured point (${\\cal M}_{\\rm dark}/{\\cal M}_{\\rm luminous}\\approx19$), making \\objectname[]{NGC~2915} one of the darkest disk galaxies known \\citep{mcbf96}. The core of the dark matter halo also appears unusually dense and compact. In fact, the stellar content can be neglected at all radii without greatly affecting the fit. The optical properties of \\objectname[]{NGC~2915}, discussed in more detail by \\citet{mmc94}, are thus unimportant for this work. We will rather focus on its unique \\ion{H}{1} properties, described more fully in \\citet{mcbf96}.  \\citet{bfpm99} discussed the origin of the \\ion{H}{1} bar and spiral pattern in \\objectname[]{NGC~2915}. They argue for a common, slow pattern speed $\\Omega_p=8.0\\pm2.4$~km~s$^{-1}$~kpc$^{-1}$, yielding a corotation radius to bar semi-length ratio $r_{\\rm cr}/r_b\\ga1.7$ ($r_b=180\\arcsec=4.6$~kpc). This lower limit is reached for $\\Omega_p=10.4$~km~s$^{-1}$~kpc$^{-1}$, while the limited extent of the circular velocity curve prevents the derivation of a proper upper limit. Although \\citeauthor{bfpm99}'s (\\citeyear{bfpm99}) measurement using \\citet{tw84a} method is plagued by uncertainties, it is consistent with the idea that bars in dense halos should be slow \\citep[e.g.][]{ds98}. Furthermore, the low (luminous) disk surface density and the locations of the pseudo-rings make it unlikely that swing amplification \\citep{t81} or bar-driving \\citep{s81,s84,brsbc94} can explain the bar and spiral patterns. Because \\objectname[]{NGC~2915} is isolated, gravitational interactions cannot be invoked as an exciting mechanism either \\citep{n87}. \\citet{bfpm99} proposed two alternatives: i) the \\ion{H}{1} disk is embedded in a massive and extended triaxial dark halo with a rotating figure, forcing the bar and spiral pattern at a certain frequency, or ii) (some) dark matter is distributed in a disk closely following the \\ion{H}{1} distribution, rendering the disk gravitationally unstable. Although both solutions lack direct observational support (but see \\citealt*{pcm94} for the latter), they offer a way forward.  The purpose of this paper is to investigate whether either (or both) of these mechanisms works in practice. The rotating halo shall be modeled by external forcing from a rotating triaxial potential (which could be due to a rotating triaxial halo, bulge, central bar, or a mixture of those), and the massive disk will be modeled by simply scaling up the gaseous surface density of the disk. We test the models by running a simple two-dimensional (hereafter 2D) hydrodynamic code, exploring the available parameter space, and comparing with the observations. In \\S~\\ref{sec:method}, we first present the relevant parameters for both sets of models and describe suitable observational constraints. A brief description of the code is provided in \\S~\\ref{sec:hydro_code}. The results for the external perturber models are presented and discussed in \\S~\\ref{sec:triaxial_halos}, while those for the massive disk models are described in \\S~\\ref{sec:heavy_disks}. We summarize all results and conclude briefly in \\S~\\ref{sec:conclusions}. ",
        "conclusions": "} \\nopagebreak We have performed 2D hydrodynamic simulations of the galaxy \\objectname[]{NGC~2915} in order to find a plausible explanation for the spiral structure observed in its extended \\ion{H}{1} disk. We have explored two main hypotheses: i) the observed bar/spiral structure is excited by an external driver, which we assume to be either a rotating bar or triaxial halo, and ii) the observed structure is due to a gravitational instability in the disk, which would contain an important unseen component and be much more massive than currently inferred from \\ion{H}{1} observations. In the first case, the free parameters are the external driver pattern frequency and strength (bar mass or halo triaxiality). Our results constrain the pattern frequency to $\\Omega_p\\approx5.5-6.5$~km~s$^{-1}$~kpc$^{-1}$ and the bar mass and halo axis ratio to $M_{\\rm bar}\\approx5-7\\times10^9$~\\msun\\ and $b/a\\approx0.83-0.88$, respectively. However, based on current structure formation simulations, it is unlikely that a triaxial halo would have such a fast figure rotation (the spin parameter $\\lambda$ is too large), while the mass required for the bar is at least an order of magnitude larger than the optical mass of the central galaxy. All our external driver models also fail to reproduce the small spiral arms pitch angle observed in the outer disk. The external perturber model is thus implausible and we do not favor it. In the heavy disk case, the only free parameter is the disk surface density scaling ratio. We found that a (corrected) scaling ratio of about $10$ leads to a spiral instability adequately accounting for the disk spiral structure and non-circular motions: the spiral pitch angle and arm-interarm contrast are accurately reproduced. The scaling ratio is well constrained by the simulations and the main uncertainty arises from the correction factor of the perturbed potential (to account for finite thickness). A scaling ratio of $10$ yields roughly the same amount of disk matter and halo dark matter over the disk radius while the halo dominates further out. This is consistent with recent $\\Lambda$CDM structure formation simulations. Furthermore, given that the heavy disk model reproduces the perturbed velocities and leads to a redistribution of angular momentum, an external triaxial potential is not required to explain the fueling of the central BCD starburst. Even for such a heavy disk, \\citeauthor{k89}'s (\\citeyear{k89}) star formation criterion is valid and accounts for the absence of star formation in the disk of \\objectname[]{NGC~2915}, provided it reads as $Q>Q_{\\rm crit}$, where $Q_{\\rm crit}$ is the critical value of the Toomre parameter for the onset of Jeans instability. As the disk of \\objectname[]{NGC~2915} does not lie in the equatorial plane of a more massive, thicker stellar disk, the expression of $Q_{\\rm crit}$ is different from that of standard galactic gaseous disks.  Throughout this work, we have neglected a possible warp in \\objectname[]{NGC~2915}, we have assumed that the observed velocity dispersion could be adequately modeled as the sound speed of an isothermal gas, and we have not considered the small-scale clumpy structure present at the beam scale in the \\ion{H}{1} observations. We have also assumed, respectively, a unique external driver pattern frequency and a uniform disk surface density scaling ratio, in order to keep the dimensionality of our parameter space as small as possible. As such, our study is exploratory in nature. This might in turn lead to too crude a description of \\objectname[]{NGC~2915}, and relaxing one or several of these assumptions may be required to improve the fit. For our heavy disk models, we also assumed a non-responsive spherically symmetric halo, whereas the disk mass is comparable to the halo mass within its radius. Considering a live halo thus appears a priority for any further similar modeling. It should in particular improve the fit to the observed central bar-like structure, where our current model fails.  Although we favor the heavy disk model over an external perturber, both models fail to reproduce satisfactorily the gas response in the center. Perturber models can loosely reproduce the bar-like feature observed but are unable to account for the kinematic twists, while heavy disk models simply fail to generate bar-like structures. The complexity of the central dynamics is in any case probably much greater than our models allow, and it is likely for instance that a different forcing frequency has to be considered for the central regions. In addition, the heavy disk models do not include any interaction or feedback from the stars, which could initiate the bar-like response. Despite a poor description of the central regions, we emphasize that a satisfactory description of the outer disk ($4.6<r<15$~kpc) is obtained, simultaneously in terms of perturbed velocity and perturbed density, with a very restricted set of adjustable parameters ($2$ in the forced runs, $1$ for the heavy disks). Furthermore, only the heavy disk models properly account for the observed small pitch angle of the spiral arms. This strong constraint thus favors this latter set of models.  If one adopts the massive disk interpretation, \\objectname[]{NGC~2915} appears as an aborted optical grand-design spiral because its gas layer has not reached the critical surface density for star formation. The spiral pattern is nevertheless present because the surface density is {\\em just} under this threshold. We recall however that there is a degeneracy between the surface density and velocity dispersion of the dark disk material. A long-standing and possibly related issue is how the \\ion{H}{1} velocity dispersion can remain at $\\approx8$~km~s$^{-1}$ in the outer disk, without any obvious thermal or mechanical energy source. In particular, the question whether the observed spiral density wave could constitute a sufficient heating source for the disk material should be addressed (whether due to an external driver or to a gravitationally unstable massive disk).  It is interesting to note that if, in any isolated object, the increased disk surface density required to explain the existence of a spiral pattern were to over-predict the observed circular velocity curve, then disk dark matter (exclusively) could be ruled out as the source of that pattern, leaving a rotating triaxial halo as the most likely source. Undisturbed LSB galaxies which are not dark matter dominated and show a clear spiral structure would be ideal candidates for such a test. However, it is not clear whether such objects exist. Spiral-based scaling factors in LSB galaxies generally produce maximum disk velocities that are high but still within the maximum allowed by the circular velocity curve \\citep[e.g.][]{f02}."
    },
    "0212/astro-ph0212080_arXiv.txt": {
        "abstract": "The majority of Galactic high-energy $\\gamma$-ray sources continue to elude identification. Currently, we have a handful of firm pulsar identifications, one of which is radio quiet, and a few marginal detections, including one millisecond pulsar. Recently, both blind searches of $EGRET$ error boxes and targeted searches of X-ray counterpart candidates have had some success in finding new pulsars. I review these results, and discuss our current program of searching mid-Galactic latitude $EGRET$ error boxes using the Parkes multi-beam system. ",
        "introduction": "The $EGRET$ $\\gamma$-ray telescope detected around 300 sources at energies above 100 MeV, the majority of which remain unidentified (Hartman et al. 1999). Large error boxes on the order of $1^{\\circ}$ across make unambiguous identifications on the basis of source position difficult. Timing observations allow definitive identifications for sources showing aperiodic (such as active galactic nuclei) or periodic (pulsars) variability. Between 6-8 pulsars have been positively identified in the $EGRET$ data by epoch folding on the known period determined from radio or, in the singular case of Geminga, X-ray data (Thompson et al 1999, Kaspi et al. 2000 and references therein). Several other energetic pulsars were found to be coincident with unidentified $EGRET$ sources, but folding the sparse $\\gamma$-ray data using contemporaneous radio ephemerides did not yield significant detections (Fierro, 1995). Although $GLAST$, when it flies in 2006, may be able to discover new pulsars  just by searching the $\\gamma$-ray data, it is likely that a priori knowledge of a pulse period will in many cases prove crucial in determining which $\\gamma$-ray sources are pulsars. ",
        "conclusions": ""
    },
    "0212/astro-ph0212563_arXiv.txt": {
        "abstract": "We investigate dynamical response on size and velocity dispersion to mass loss by supernovae in formation of two-component spherical galaxies composed of baryon and dark matter.  Three-dimensional deprojected de Vaucouleurs-like and exponential-like profiles for baryon, embedded in truncated singular isothermal and homogeneous profiles for dark matter, are considered.  As a more realistic case, we also consider a dark matter profile proposed by Navarro, Frenk \\& White.  For simplicity we assume that dark matter distribution is not affected by mass loss and that the change of baryonic matter distribution is homologous.  We found that the degree of the response depends on the fraction of dark matter in the region where baryon is distributed, so that dwarf spheroidal galaxies would be affected even in a dark halo if they are formed by galaxy mergers in the envelope of the dark halo.  Our results suggest that this scenario, combined with dynamical response, would make not only the observed trends but the dispersed characteristics of dwarf spheroidals. ",
        "introduction": "Elliptical galaxies have provided important information on the formation and evolution of galaxies.  So far some scaling relations among photometric properties and global structural parameters have been observed, such as the colour--magnitude relation \\citep[e.g.,][]{b59}, the velocity dispersion--magnitude relation \\citep{fj76}, and the surface brightness--size relation \\citep{k77, kow83}.  These relations are regarded as clues to our understanding of galaxy formation.  It is widely believed that heating up and sweeping out the galactic gas by multiple supernova (SN) explosions play a key role on the formation of dwarf elliptical/spheroidal galaxies.  In a traditional framework of monolithic cloud collapse scenario for formation of ellipticals, such multiple SNe cause a galactic wind to halt subsequent star formation and chemical enrichment \\citep{l69, i77, s79, ay86, ay87, ka97}.  Since a large amount of gas is expelled by the galactic wind, the self-gravitating system expands after blowing up the wind \\citep[e.g.,][]{h80, m83, v86}.  \\citet{ya87} found that many observed properties of ellipticals can be reproduced by the galactic wind model with the evolutionary population synthesis technique.  They noticed that theoretical metallicity of galaxies, to be compared with observation, should be an average of metallicities of constituent stars weighted by their luminosities.  On the other hand, \\citet{ds86} considered a model in which galaxies are embedded in dominant dark haloes.  They claimed that structural properties of luminous matter would not be affected by the gas removal if dark matter dominates the gravitational potential. Their model can also reproduce the observed relations, although they did not take into account the evolutionary population synthesis.  Recently, \\citet{z02} argued the dynamical limit on occurrence of galactic wind given by the existence of sufficiently dense dark halo and the absence of intergalactic/escaping globular clusters, and concluded that only about 4 per cent of the total mass of the galaxy should be lost by the wind.  In contrast to the monolithic cloud collapse scenario, recent studies on the cosmological structure formation suggest the hierarchical clustering scenario based on a cold dark matter (CDM) model.  In this scenario, the CDM dominates the Universe gravitationally and objects or dark haloes continuously cluster and merge together, so that larger objects form via mergers of smaller objects.  Thus, simple galactic wind models must be modified in order to be consistent with the cosmological structure formation.  Based on the hierarchical clustering scenario, semi-analytic models (SAMs) of galaxy formation have been developed \\citep[e.g.,][]{kwg93, c94, c00, sp99, ngs99, n01, n02}.  These models well reproduce many statistical observables of galaxies such as luminosity function, colour distribution, and gas fraction. \\citet{kc98} and \\citet{ng01} have shown that their SAMs taking into account metal enrichment can reproduce the observed colour--magnitude relation of elliptical galaxies in clusters of galaxies and that this relation reflects the mass-metallicity relation, as pointed out by using the galactic wind model \\citep{ka97}.  Thus, as a next step, it is valuable to investigate not only the photometric properties but structural properties such as size and velocity dispersion.  Moreover, \\citet{on01a, on01b} found that radial gradients of fractional number, colour, and star formation rate of ellipticals in clusters are also reproduced by using the SAM associated with an $N$-body simulation.  Because a large amount of mass loss might occur during the formation of ellipticals, in both scenarios of monolithic cloud collapse and hierarchical clustering, it is important to construct a formalism on the dynamical response in two-component galaxies consisting of baryon and dark matter.  While self-consistent equilibrium models of such two-component elliptical galaxies have been analysed by \\citet{ys87} and \\citet{cp92}, we focus on the dynamical response of baryon within a dark matter halo.  In this paper we consider simple models of two-component galaxies in which the baryon having either de Vaucouleurs or exponential profile is embedded in the dark matter having either isothermal or homogeneous profile.  Moreover, we also consider the so-called NFW profile of dark matter that is recently suggested by $N$-body simulations \\citep{nfw}.  This paper is outlined as follows.  In section 2, general formulation for the dynamical response on size and velocity dispersion to mass loss is presented.  In section 3 the dynamical response for particular choices of density distributions mentioned above is investigated.  Section 4 summarizes the results of this paper.  The details of models and derivations are given in Appendices. ",
        "conclusions": "We analyzed the dynamical response on structural parameters of elliptical galaxies to mass loss.  While in the previous analyses only two limiting cases of pure self-gravitating baryonic system and dark matter-dominated system were considered \\citep[e.g.,][]{ya87, ds86}, we extended the formalism to the intermediate case between the two. Three-dimensional density distributions of baryonic matter considered in this paper are the Jaffe model that approximates the de Vaucouleurs profile and a model that approximates the exponential profile in two dimensions.  As for the density distribution of dark matter, a truncated isothermal sphere and a homogeneous sphere are considered.  As a more realistic case, we also considered the NFW dark matter halo which surrounds the baryonic component having de Vaucouleurs profile.  Thus we found, from five combinations of baryon and dark matter, that the response strongly depends on the central concentration of dark matter, reflecting the fraction of baryon relative to dark matter in the central region.  In contrast to the simple assumption by \\citet{ds86} that structural parameters are not affected by mass loss, size and velocity dispersion can change even in dark matter, if baryon has sufficiently shrunk and become dense compared to dark matter at the onset of mass loss.  In this paper we considered both adiabatic and instantaneous gas removal.  The adiabatic removal is justified if the timescale of gas removal is sufficiently slow compared to the dynamical timescale of the system.  On the other hand, in the case of instantaneous gas removal, the effects of gas removal are stronger.  For example, when more than a half of the mass is removed instantaneously, the pure self-gravitating system becomes unbound.  Thus, our results for adiabatic removal place a lower limit on the dynamical response to mass loss.  In order to check the opposite case, we also derived the results for instantaneous gas removal.  As expected, the changes of size and velocity dispersion are larger than those for adiabatic gas removal.  Such drastic changes may not validate our simple assumptions of homologous expansion of baryon and static distribution of dark matter.  So far some scaling relations such as the velocity dispersion-magnitude relation have been discussed considering the dynamical response to mass loss in the context of monolithic cloud collapse scenario for formation of elliptical galaxies,  However, since the resuting structural change of baryonic component is found to depend significantly on the density distribution of dark matter, such scaling relations in the hierarchical clustering scenario should necessarily be reconsidered.  Our results suggest that the dispersed properties of dwarf/compact ellipticals would be well reproduced by a variety of their formation histories and environment, when taking into account mass-dependent gas removal.  In the hierarchical clustering scenario, galaxy formation processes are rather complicated.  There are many physical processes such as radiative gas cooling and galaxy merger as well as star formation and SN feedback. It should be noted that the physical mechanism of galaxy merger would be dynamical friction and/or random collision, while the mechanism of dark halo merger is growth of density fluctuation.  In order to investigate the observed properties of galaxies, we need to fully incorporate the effects of dynamical response into realistic models of galaxy formation like our SAMs.  We will discuss various observables of ellipticals in more realistic situations in a forthcoming paper."
    },
    "0212/astro-ph0212425_arXiv.txt": {
        "abstract": "The coming flood of CMB polarization experiments, spurred by the recent detection of CMB polarization by DASI and WMAP, will be confronted by many new analysis tasks specific to polarization.  For the analysis of CMB polarization data sets, the devil is truly in the details. With this in mind, we present the details of the data analysis for the POLAR experiment, which recently led to the tightest upper limits on the polarization of the Cosmic Microwave Background Radiation at large angular scales. We discuss the data selection process, mapmaking algorithms, offset removal, and the likelihood analysis which were used to find upper limits on the polarization. Stated using the modern convention for reporting CMB Stokes parameters, these limits are 5.0 \\uK\\ on both $E$-type and $B$-type polarization at 95\\% confidence. Finally, we discuss simulations used to test our analysis techniques and to probe the fundamental limitations of the experiment. ",
        "introduction": "The detection of polarization in the Cosmic Microwave Background (CMB) has been a long sought goal for cosmology.  CMB polarization was recently detected at sub-degree angular scales by the DASI instrument, a ground-based interferometer \\cite{dasi02a}, and the temperature-polarization correlation has now been detected at larger angular scales by the WMAP satellite \\cite{kogut03}. While information in the CMB polarization at degree and sub-degree angular scales further confirms the standard cosmological model, polarization information at larger angular scales has the potential to provide additional information regarding the formation and evolution of the universe.  The process of reionization leaves a characteristic signature on CMB polarization at large angular scales which can be used as a means to determine the epoch of reionization \\cite{matias97,hu2000,kap02}. The power spectrum of CMB polarization at angular scales greater than a degree or so is sensitive to inflationary model parameters such as the inflaton potential and the energy scale of inflation, as well as primordial gravitational waves \\cite{zs97,kks,kin98,knox02,kck03}.  These exciting rewards, taken together with increasingly sensitive receiver technologies, have set the stage for a host of new CMB polarization experiments.  These experiments will be faced with new and more challenging analysis tasks than for the simpler case of anisotropy, and it is with that in mind that we now set out to report the details of the data analysis of the \\polar\\ experiment.  \\polar\\ (Polarization Observations of Large Angular Regions) was designed to measure the polarization of the CMB on large angular scales, in the \\ka\\ frequency band from $26-36$ GHz. This HEMT-based correlation polarimeter operated for a single season in the Spring of 2000 near Madison, Wisconsin. The data from this single season led to simultaneous upper limits on E and B-type CMB polarization, results which were initially presented in \\cite{kea01} (hereafter K01). The details of the \\polar\\ instrument and its operation were later described in \\cite{kea02} (hereafter K02). In this paper we present the details of the data selection and analysis techniques used to arrive at the results in K01.  We also discuss the results of a recent cross-correlation analysis of the POLAR data with anisotropy data from the COBE-DMR experiment, described fully in \\cite{adoc02}.  The rest of this paper is organized as follows.  We first briefly discuss the two conventions used to quantify CMB polarization. We next review the basic properties of the POLAR instrument in \\sct{s:instr}. In \\sct{s:cuts}, we discuss the data selection procedures that were used to remove large amounts of contaminated data. \\sct{s:mapmaking} describes the mapmaking algorithms used to transform the raw data into sky maps of $Q$ and $U$, and provides a full pipeline simulation in order to test the algorithms. \\sct{s:likes} presents the likelihood analysis used to arrive at the upper limits on CMB polarization, as well as the evaluation of the polarization power spectra and some commentary on the lack of substantial foreground contamination. Finally, in \\sct{s:discussion}, we discuss the limitations of our experiment that could be improved upon in future projects. ",
        "conclusions": "\\label{s:discussion} To summarize our results thus far, using the \\polar\\ data we have placed limits on CMB polarization power spectra in the context of both a flat band-power model and a concordance model, and we have discussed the lack of substantial synchrotron contamination in our scan region.  We showed in \\sct{s:fbp} that the offset removal technique only degraded our limits by $\\sim 30\\%$, which is not so bad considering that without it, no limits would have been possible.  If we had integrated for a much longer time and/or our detectors had been much more sensitive, how well could we have actually done in determining $E$ and $B$ on large angular scales given our scan strategy? We attempted to answer this question via an additional simulation. We used CMBFast \\cite{cmbfast} to generate power spectra for a concordance model universe with no $B$-modes (\\ie, only scalar perturbations) and an optical depth to reionization of $\\tau=0.2$; ideally we should recover only $E$ modes from the analysis.  We used the HEALPix package \\cite{healpix} to generate twenty sky realizations of $Q$ and $U$, and observed each with a nearly perfect \\polar, one that had a sensitivity of 0.01 \\uK\\ for each of 180 2\\deg\\ wide pixels in our $\\delta=43\\deg$ RA strip. We then performed our likelihood analysis on these map realizations limited to our scan region, and found that we recovered roughly equal amounts of $E$ and $B$ power. This is consistent with the window functions generated in our quadratic estimator analysis of \\sct{qe}.  On a positive note, we did recover the correct amount of total $E$ and $B$ power; that is, the recovered $C_\\l^E + C_\\l^B$ equaled the initial $C_\\l^E$ placed in the maps.  Ultimately, \\polar\\ was limited both by its instrument sensitivity and atmospheric limitations to integration time.  The atmosphere was also most likely responsible for the time-varying polarization offset. Finally, our one-dimensional scan strategy prevented discrimination between E-modes and B-modes. Future experiments that effectively deal with these issues may well be able to glean the CMB polarization signals on large angular scales, and hence uncover a wealth of new information with which to increase our understanding of inflation and the details of the Big Bang."
    },
    "0212/astro-ph0212343_arXiv.txt": {
        "abstract": "Statistical properties of gas absorption in high redshift quasars such as power spectrum and bispectrum allow one to determine the evolution of structure over the redshift range $2<z<4$. Sloan Digital Sky Survey (SDSS) will measure around 10,000 quasar spectra in this redshift range and will allow one to determine the growth factor with a few percent accuracy. This allows one to extend the studies of dark energy to high redshift and determine the presence of dark energy if $\\Omega_{\\rm de}>0.1-0.2$ at $z>2$. In combination with low redshift studies one can place useful limits on the time evolution of the equation of state. ",
        "introduction": "The study of the Ly-$\\alpha$ forest has been revolutionized in recent years by the high resolution measurements using Keck HIRES spectrograph and by the development of theoretical understanding using hydrodynamical simulations and analytical models. The picture that has emerged from these studies is one where the neutral gas responsible for the absorption is in a relatively low density, smooth enviroment, which implies a simple connection between the gas and the underlying dark matter. The neutral fraction of the gas is determined by the interplay between the recombination rate (which depends on the temperature of the gas) and ionization caused by ultraviolet photons, the so called UV background. Photoionization heating and expansion cooling cause the gas density and temperature to be tightly related, except where mild shocks heat up the gas. This leads to a tight relation between the absorption flux and the gas density. Finally, the gas density is closely related to the dark matter density on large scales, while on small scales the effects of thermal broadening and Jeans smoothing have to be included. In the simplest picture described here all of the physics ingedients are known and can be modelled. In practice, the nonlinear physics requires the use of hydrodynamic simulations with sufficient dynamic range which are only now becoming available.  In the past couple of years cosmological observations such as redshift-luminosity relation from supernovae and position of the acoustic peaks in cosmic microwave background (CMB) combined with local matter density estimates have revealed the presence of another component in the universe, the so called dark energy. This component has negative pressure and the recent constraints indicate that $w=p/\\rho \\sim -1$. However, at the moment direct constraints are still allowing for a significant range in $w$, specially if it is allowed to vary in time. Many models have been proposed where $w$ either increases or decreases with redshift (some of these are discussed in these proceedings). Theoretical studies have been performed on how to extract this evolution using low redshift probes. While statistical power for some of these planned or proposed experiments, such as SNAP using supernovae, LSST using weak lensing or optical, X-ray and Sunyaev-Zeldovich telescopes counting clusters, is truly impressive, the control of systematics is less well understood. For this reason it is important to have as many independent tests as possible.  One possibility that has not been discussed much so far is using Ly-$\\alpha$ forest as a probe of dark energy. Since direct observations of quasars from the ground restrict one to the optical wavelengths ($\\lambda>3600\\AA$) one can only observe Ly-$\\alpha$ forest for $z>2$. If $w=-1$ independent of the redshift then if dark energy to dark matter density ratio $\\Omega_{\\rm de}/\\Omega_m=2$ today the ratio will be below 0.1 at $z>2$. In this case the growth factor will scale almost linearly with the scale factor just as in an Einstein-de Sitter universe and one cannot detect the presence of dark energy. On the other hand, if $w>-1$ either today or at a somewhat higher redshift then the decline of the dark energy fraction relative to dark matter is slowed down and dark energy can be dynamically important even at $z>2$. In this case there will be deviations in the growth factor from the expected scaling that can be detected using the forest observations.  So far the only tests of the dark energy proposed have been using either the growth factor or the redshift luminosity distance. If one wishes to extract $w(z)$ then both of these involve a double integral over this quantity and degeneracies arise. A more direct and still in principle observable way is to measure the Hubble parameter $H(z)$, which is related to $w(z)$ through a single integral. One way to measure it is to have a characteristic feature which is fixed in comoving coordinates and observed in redshift space. The relation between redshift space and comoving space is the Hubble parameter itself and so observing the feature as a function of redshift provides $H(z)$ directly. The problem of course is that there are no characteristic features imprinted in large scale structure, since the distribution of structure in the universe is stochastic in nature. One must therefore look for a characteristic scale in correlations between structures. In principle such features could be provided by baryonic oscillations imprinted in the matter power spectrum, but in practice this is a weak effect limited to very large scales and so cannot be made very precise. One is thus left with the variations in the correlation function slope as a function of scale. It is well known that the power spectrum slope varies from $n\\sim 1$ on large scales to $n \\sim -3$ on small scales independent of redshift in CDM models. Hence the scale at which the slope takes a specific value can be viewed as a standard ruler and can be traced with redshift. If this slope is measured in redshift space, as is the case for either galaxy clustering along the line of sight or Ly-$\\alpha$ forest correlations, then one is measuring directly $H(z)$. To be able to do this one has to detect the curvature in the slope over the dynamic range of observations. This is challenging, since the dynamic range is narrow and the error on the slope will be large. We will show below that such a detection should be possible with the current sample, but is not expected to improve the constraints on $w(z)$ significantly. ",
        "conclusions": ""
    },
    "0212/astro-ph0212496_arXiv.txt": {
        "abstract": "{ The results of astrometric and photometric investigations of the open cluster NGC$\\,$7243 are presented. Proper motions of 2165 stars with  root-mean-square error of $1.1\\;{\\rm mas\\,yr^{-1}}$ were obtained by means of PDS scanning of astrometric plates covering the time interval of 97 years.         A total of 211 cluster members down to $V=15.5\\;$mag have been identified. $V$ and $B$ magnitudes have been determined for 2118 and 2110 stars respectively. Estimations of mass ($348 {M_\\odot} \\le M_{TOT} \\le 522 {M_\\odot}$), age $(t=2.5\\cdot 10^8\\,$yr), distance $(r=698\\,$pc) and  reddening $(E_{B-V}=0.24)$ of the cluster NGC$\\,$7243 have been made. ",
        "introduction": "The open cluster NGC\\,7243 is located in the Lacerta constellation. It has the following coordinates: $\\ell=98\\fdg 9$, $b=-5\\fdg 6$, $\\alpha=22^{\\rm h} 15^{\\rm m}$, $\\delta=+49\\degr 53\\arcmin$ (2000.0). In the Lyng\\aa\\, (1987)                  % catalogue it is classified as Tr\\\"umpler (1930) class $\\,$II$\\,$2m with % 40 cluster members and an angular diameter of $21\\arcmin$. Early investigations of this cluster led Raab (1922)         % to the conclusion that, in fact, it consisted of two clusters, but a lack of sufficient observational data made it impossible to check this statement at that epoch. Zug (1933)           % was the first investigator to study the interstellar absorption in the cluster area. He denoted the brightest stars and these numbers are still used in some papers. Relative proper motions of 814 stars in the cluster area were for the first time obtained by Lengauer (1937)            % from  plates taken with the Normal Astrograph of the Pulkovo Observatory with a 33 years mean epoch difference. As a result, 32 cluster members were selected. Later Rahmatov and Muminov (1985)                        % determined proper motions of 1287 stars in the region of NGC\\,7243 using the plates of the Normal Astrograph in Tashkent (Uzbekistan) as the first epoch plates. Due to intensive urban light pollution, the second epoch plates had to be taken with the 40\\,cm Zeiss Astrograph of the Kitab Latitude Station (Uzbekistan). In the circle of $40\\arcmin$ radius, 28 possible cluster members with a membership probability of more than $50\\%$ were selected. van Schewick (1957) using only one pair of plates          % confirmed the existence of  cluster members, previously found by Lengauer, but did not add any one new. Unfortunately, none of the above mentioned  investigations, for one reason or another (absence of photometric data, small epoch difference, small number of plate pairs, etc.) permitted  to select all probable  members of the cluster NGC\\,7243 in the observed magnitude range and to estimate membership probabilities.  The Main Astronomical Observatory at Pulkovo has 19 first epoch plates (obtained from year 1897 to 1949). This allows to solve the astrometric part of the problem with a high degree of confidence. For a reliable segregation of the cluster members it is necessary to use photometric data based on the international photometric system.  Photometric plates of the NGC\\,7243 region in $B$ and $V$ passbands  were obtained with the Schmidt Telescope (80/120/240) of the Latvian Radioastrophysical Observatory. ",
        "conclusions": "We have carried out astrometric and photometric study of the young  open cluster  NGC\\,7243  poorly investigated before. The main  cluster parameters obtained in this investigation are given in Table~7. The complete catalogue of 2623 stars in the region of the NGC$\\,$7243 containing positions, proper motions and their individual errors, $B$ and $B-V$ magnitudes and individual membership probabilities (Table 8) is available in electronic form at the CDS via anonymous ftp 130.79.128.5.  \\bigskip  \\begin{table*} \\caption{NGC$\\,$7243 summary data} \\begin{center} \\begin{tabular}{ll} \\bigskip  Number of members              &    $211$\\\\ Limiting magnitude             &    $V \\sim 15.5\\;{\\rm mag}$\\\\ Magnitude of the brightest star &   $V = 8.43\\;{\\rm mag} $\\\\ Earliest spectral type         & B5$\\,$III \\\\ MS turn-off point              & $(B-V)_0 = - 0.11\\;{\\rm mag}$\\\\ Age                            & $t=(2.5\\pm 0.5)\\cdot 10^8\\,$yr\\\\ Total cluster mass    & $348M_{\\odot} \\le M_{TOT}\\le 522M_\\odot$ \\\\ Average reddening     & $E_{B-V}=0.24\\;{\\rm mag}$\\\\ Distance              &  $r =  {698^{\\,+130}}\\!\\!\\!\\!\\!\\!\\!\\!\\!\\!\\!_{-109}\\,$pc \\\\ Mass function slope            & $-2.37 \\pm 0.27$  \\\\ \\end{tabular} \\end{center} \\end{table*}  \\begin{table*} \\begin{center} \\caption[]{Complete catalogue of 2623 stars in the region of the NGC\\,7243 (example of the Table~8 only available in electronic form at the CDS)} \\begin{tabular}{rccccccccccccccc} \\hline This &RS&Len-& $\\alpha$   &   $\\delta$   & Pair & $\\mu_x$& $\\epsilon_x$ & $\\mu_y$&$\\epsilon_y$& $V$&$B-V$& $P$&  $X$&$Y$& Notes \\\\ work  &  &gauer & \\multicolumn{2}{c}{2000.0} &      & \\multicolumn{4}{c}{${\\rm mas\\, yr^{-1}}$}          &    &    &  \\%   &   mm    & mm &\\\\ \\hline 1 & & & 22 10 42.86 &  49 17 39.2  &   9 &    1.7 &  1.0 &    0.9 &   1.2 &   11.45: &    1.20: &   97 & -43.67 &  -36.14 & \\\\ 2 & & &22 10 46.88 &  49 16 44.2  &   9 &    1.0 &  1.7 &   -2.3 &   1.6 &   13.40  &    0.77  &   62 & -43.02 &  -37.07 &\\\\ 3 & & &22 10 48.33 &  49 19 02.6  &   9 &    8.4 &  1.0 &   -9.8 &   0.8 &   11.78  &    0.74  &    0 & -42.75 &  -34.75&\\\\ 4 & & &22 10 49.61 &  49 14 59.1  &   9 &   -2.1 &  1.2 &   -0.7 &   0.8 &   11.98: &    0.40: &    1 & -42.60 &  -38.85& \\\\ ...&...&...&...       &...           &...  &...     &...   &...     &...    &...       &...       &...   &...     &...     & \\\\ 289 & &717& 22 15 43.11&49 26 57.4  &  10 &     3.7 &  1.0&    1.9 &   0.7 &    8.71: &    0.14:&    89 &   5.64 &  -27.13& HD\\,211418 \\\\ ... &... &... &... &... &... &... &... &... &... &... &... &... &... &...  \\\\ 2623&&&  22 10 34.29  &  50 29 02.3  &  10  &    0.9  &  0.7  &   -0.9 &   0.8  &  12.40: &    0.57: &   91 &  -43.93 &   35.79 \\\\ \\hline &&&&&&&&&&\\\\ \\multicolumn{14}{l}{\\small Column (1) star number in our catalogue,} \\\\ \\multicolumn{14}{l}{\\small Column (2) the reference stars are marked with *,} \\\\ \\multicolumn{14}{l}{\\small Column (3) star number according to Lengauer,} \\\\ \\multicolumn{14}{l}{\\small Columns (4) and (5)  equatorial coordinates at 2000.0,}\\\\ \\multicolumn{14}{l}{\\small Column (6)  number of used astroplate pairs,}\\\\ \\multicolumn{14}{l}{\\small Columns (7) and (8)  proper motion $\\mu_x$ and its error  $\\epsilon_x$,}\\\\ \\multicolumn{14}{l}{\\small Columns (9) and (10) proper motion $\\mu_y$ and its error  $\\epsilon_y$,}\\\\ \\multicolumn{14}{l}{\\small Column (11)  $V$ magnitude,}\\\\ \\multicolumn{14}{l}{\\small Column (12) $B-V$ magnitude,}\\\\ \\multicolumn{14}{l}{\\small Column (13)  value of membership probability $P$,}\\\\ \\multicolumn{14}{l}{\\small Column (14) and (15) rectangular coordinates (relative to the star 898) in the scale of the Normal Astrograph,}\\\\ \\multicolumn{14}{l}{\\small Column (16) notes}\\\\ \\end{tabular} \\end{center} \\end{table*}"
    },
    "0212/astro-ph0212175_arXiv.txt": {
        "abstract": "We present near-infrared imaging and spectroscopic observations of two FR~II high-redshift radio galaxies (HzRGs), 4C~40.36 ($z=2.3$) and 4C~39.37 ($z=3.2$), obtained with the Hubble, Keck, and Hale Telescopes.  High resolution images were taken with filters both in and out of strong emission lines, and together with the spectroscopic data, the properties of the line and continuum emissions were carefully analyzed.  Our analysis of 4C~40.36 and 4C~39.37 shows that strong emission lines (e.g., \\oiii\\ 5007 \\AA\\ and \\hanii) contribute to the broad-band fluxes much more significantly than previously estimated (80\\% vs. 20--40\\%), and that when the continuum sources are imaged through line-free filters, they show an extremely compact morphology with a high surface brightness.  If we use the $R^{1/4}$-law parametrization, their effective radii ($r_{e}$) are only 2--3 $h_{50}^{-1}$ kpc while their restframe $B$-band surface brightnesses at $r_{e}$ are $I_{e}(B) \\sim$ 18 mag/$\\sq$\\arcsec. Compared with $z \\sim 1$ 3CR radio galaxies, the former is $\\times$ 3--5 smaller, while the latter is 1--1.5 mag brighter than what is predicted from the $I_{e}(B)-r_{e}$ correlation.  Although exponential profiles produce equally good fits for 4C~40.36 and 4C~39.37, this clearly indicates that with respect to the $z\\sim1$ 3CR radio galaxies, the light distribution of these two HzRGs is much more centrally concentrated.  Spectroscopically, 4C~40.36 shows a flat ($f_{\\nu}$=const) continuum while 4C~39.37 shows a spectrum as red as that of a local giant elliptical galaxy.  Although this difference may be explained in terms of a varying degree of star formation, the similarities of their surface brightness profiles and the submillimeter detection of 4C~39.37 might suggest that the intrinsic spectra is equally blue (young stars or an AGN), and that the difference is the amount of reddening. ",
        "introduction": "Because of their high luminosities, radio galaxies are extremely useful for studying the process of galaxy formation and evolution at high redshift.  Although their substantial luminosities---the very property which makes them so useful---may make them quite atypical of galaxies, whatever is learned for high-redshift radio galaxies (HzRGs) might also be applicable to a much broader spectrum of high-redshift galaxies.  When using radio galaxies as a probe for galaxy formation/evolution at high redshift, the important observations are of the near-IR continuum emission.  Although the spectacular morphology of the luminous line-emitting regions attracts our attention, the line-emitting regions do not provide significant information on the underlying stellar populations since they are most likely excited by the central AGN \\citep{Rawlings89,Willott99}.  On the other hand, the near-IR continuum emission is thought to originate from the underlying stellar population because of the remarkable tightness and continuity seen in the $K$-band Hubble diagram ($K-z$ relation) of HzRGs as first shown by \\citet{Lilly84} (for the recent compilations of $K-z$ diagram, see \\citet{Jarvis01} and \\citet{Breuck02}).  Since the UV continuum is often contaminated by scattered AGN light \\citep[e.g.,][]{Vernet01}, the near-IR continuum is the most robust observable from which we can hope to extract information on the underlying stellar populations in HzRGs.  When deriving near-infrared continuum properties of HzRGs such as magnitudes and morphology, it is essential to remove the contribution from the luminous line-emitting regions \\citep[e.g.,][]{Eales93}. This becomes especially important when observing radio galaxies at $2 \\lesssim z \\lesssim 4$, where strong optical emission lines often contaminate $H$- and $K$-band observations.  This redshift range is of great interest since HzRGs start to show a dramatic evolution in terms of morphology \\citep{Breugel98} and submillimeter luminosity \\citep{Archibald01}, although the $K$-band (i.e., restframe visual) luminosity evolution remains relatively mild \\citep{Jarvis01}.  To study the properties of the near-infrared continuum sources in HzRGs, we have been carrying out a near-infrared imaging and spectroscopic study of HzRGs using the Keck, Hale, and Hubble Space Telescopes (HST) (\\citealt{Armus98}; \\citealt{Egami99}; Armus~et~al.~2002, in preparation).  In this paper, we present observations of two FR~II radio galaxies, 4C~40.36 at $z=2.27$ \\citep{Chambers88} and 4C~39.37 ($=$ 6C~1232$+$39) at $z=3.22$ \\citep{Rawlings90,Eales93b}.  These galaxies were chosen for HST/NICMOS observations because their restframe visual emission lines (\\oiii\\ and \\hanii\\ for 4C~40.36 and \\oiii\\ for 4C~39.37) are redshifted into the narrow-band filters available with HST/NICMOS. Together with the high-quality ground-based data from the Keck and Hale Telescopes, we are able to examine the nature of these galaxies in detail.  To facilitate the comparison with previously published results, we assume $\\Omega_{M}=1$, $\\Omega_{\\Lambda}=0$, $H_{0}=50$ $h_{50}$ km s$^{-1}$ Mpc$^{-1}$ throughout the paper.  With these parameters, 1\\arcsec\\ subtends 7.9, and 7.1 $h_{50}^{-1}$ kpc for 4C~40.36 and 4C~39.37, respectively. ",
        "conclusions": "Using the Hubble, Keck, and Hale Telescopes, we have obtained near-infrared images and spectra of two FR II HzRGs, 4C~40.36 ($z=2.27$) and 4C~39.37 ($z=3.22$).  The main conclusions are as follows: \\begin{enumerate} \\item The contributions of the \\oiii\\ 5007 \\AA\\ and \\hanii\\ lines to near-IR broad-band magnitudes are found to be $\\sim$ 80\\%, substantially larger than the previous estimates of 20--40\\% determined for the same galaxies.  \\item The continuum sources in 4C~40.36 and 4C~39.37 are extremely compact ($r_{e} \\sim 2-3$ kpc) and of high surface brightness ($I_{e}(B) \\sim 18$ mag/$\\sq$\\arcsec) compared to $z \\sim 1$ 3CR radio galaxies.  However, their restframe $B$-band luminosities are similar, showing very little luminosity evolution.  \\item The continuum spectral shapes of the two HzRGs in the restframe visual are very different: the continuum spectrum of 4C~40.36 is almost completely flat ($f_{\\nu}$=const) while the spectrum of 4C~39.37 is as red as that of a local gE galaxy.  The origin of the continuum emission is not yet clear since no obvious AGN feature (e.g., a broad line) or stellar feature (e.g., a continuum break) is seen.  \\item The difference of the continuum spectral shapes may be explained in terms of a varying degree of star formation as well as different amounts of reddening.  \\end{enumerate}"
    },
    "0212/astro-ph0212205_arXiv.txt": {
        "abstract": "We discuss the statistical mechanics of violent relaxation in stellar systems following the pioneering work of Lynden-Bell (1967). The solutions of the gravitational Vlasov-Poisson system develop finer and finer filaments so that a statistical description is appropriate to smooth out the small-scales and describe the ``coarse-grained'' dynamics. In a coarse-grained sense, the system is expected to reach an equilibrium state of a Fermi-Dirac type within a few dynamical times. We describe in detail the equilibrium phase diagram and the nature of phase transitions which occur in self-gravitating systems. Then, we introduce a small-scale parametrization of the Vlasov equation and propose a set of relaxation equations for the coarse-grained dynamics. These relaxation equations, of a generalized Fokker-Planck type, are derived from a Maximum Entropy Production Principle (MEPP). We make a link with the quasilinear theory of the Vlasov-Poisson system and derive a truncated model appropriate to collisionless systems subject to tidal forces.  With the aid of this kinetic theory, we qualitatively discuss the concept of ``incomplete relaxation'' and the limitations of Lynden-Bell's theory. ",
        "introduction": "\\label{intro}  It has long been realized that galaxies, and self-gravitating systems in general, follow a kind of organization despite the diversity of their initial conditions and their environement. This organization is illustrated by morphological classification schemes such as the Hubble sequence and by simple rules which govern the structure of individual self-gravitating systems. For example, elliptical galaxies display a quasi-universal luminosity profile described by de Vaucouleurs' $R^{1/4}$ law and most of globular clusters are well fitted by the Michie-King model \\cite{bt}. The question that naturally emerges is, what determines the particular configuration to which a self-gravitating system settles. It is possible that their present configuration crucially depends on the conditions that prevail at their birth and on the details of their evolution. However, in view of their apparent regularity, it is tempting to investigate whether their organization can be favoured by some fundamental physical principles like those of thermodynamics and statistical physics. We ask therefore if the actual states of self-gravitating systems are not simply more probable than any other possible configuration, i.e. if they cannot be considered as maximum entropy states.  This thermodynamical approach may be particularly relevant for globular clusters and elliptical galaxies which are described by a distribution function that is almost isothermal \\cite{bt}. In the case of globular clusters, the relaxation proceeds via two-body encounters and this collisional evolution is well-described by kinetic equations of a Fokker-Planck-Landau type for which a H-theorem \\footnote{The H-theorem, proved by Boltzmann for an ideal gas, states that entropy is a monotonically increasing function of time. Boltzmann's kinetic theory of gases therefore provides a direct justification of the second principle of thermodynamics. The generalization of this theorem for self-gravitating systems rests on simplifying assumptions which are difficult to rigorously justify \\cite{Kandrup}.} is available. By contrast, for elliptical galaxies, two-body encounters are completely negligible (the corresponding relaxation time $t_{{\\it coll}}$ exceeds the age of the universe by many orders of magnitude) and the galaxy dynamics is described by the {\\it Vlasov equation}, i.e. collisionless Boltzmann equation \\cite{bt}. Since the Vlasov equation rigorously conserves entropy, a relaxation towards an isothermal distribution looks at first sight relatively surprising. Yet, the inner regions of elliptical galaxies appear to be isothermal and this fact stemed as a mystery for a long time.  In a seminal paper, Lynden-Bell \\cite{lb} argued that the violently changing gravitational field of a newly formed galaxy leads to a redistribution of energies between stars and provides a mechanism analogous to a relaxation in a gas. The importance of this form of relaxation had previously been stressed by a number of authors including H\\'enon and King but Lynden-Bell showed for the first time the relevance of a statistical description. He argued that the Vlasov-Poisson system develops an intricate {\\it mixing process} in phase space associated with the heavily damped oscillations of a protogalaxy initially far from mechanical equilibrium and collapsing under its own gravity. As a result, the solutions of the Vlasov equation are not smooth but involve intermingled filaments at smaller and smaller scales. In this sense, there is no convergence towards equilibrium but rather the formation of a fractal-like structure in phase space.  However, if we introduce a macroscopic level of description and make a local average of the distribution function over the filaments, the resulting ``coarse-grained'' distribution function is smooth and is expected to reach a maximum entropy state (i.e. most mixed state) on a very short time scale of the order of the dynamical time $t_{D}$. This process is called {\\it violent relaxation} and is acknowledged to account for the regularity of elliptical galaxies or other collisionless self-gravitating systems.  Lynden-Bell predicted that the equilibrium state should be described by a Fermi-Dirac distribution function or a superposition of Fermi-Dirac distributions. Here, degeneracy is due not to quantum mechanics but to the Liouville theorem that prevents the smooth distribution function from exceeding the maximum of its initial value. In the non degenerate limit, the Fermi-Dirac distribution functions reduce to Maxwellians. The prediction of an isothermal distribution for collisionless stellar systems was considered as a triumph in the 1960's and laid the foundation of a new type of statistical mechanics. Of course, the validity of the theory is conditioned by a hypothesis of ergodicity which may not be completely fulfilled. This is the complicated problem of ``incomplete relaxation'' which limitates the power of prediction of Lynden-Bell's approach. However, as we shall see, these difficulties should not throw doubt on the importance of this statistical description. A similar relaxation process is at work in two-dimensional turbulence (described by the 2D Euler equation) and can explain the organization and maintenance of coherent vortices, such as the Great Red Spot of Jupiter, which are common features of large-scale geophysical or astrophysical flows \\cite{rs1,miller,csfluide,rr,shallow,bouchet1,bouchet2}. The mathematical relevance of this statistical description has been given by Robert \\cite{rob} introducing the concept of Young measures. The formal analogy between two-dimensional vortices and stellar systems has been discussed by Chavanis \\cite{cthese,cfloride,japon,csr}.   This paper is organized as follows. In section \\ref{VP}, we introduce the gravitational Vlasov-Poisson system and list its main properties. In section \\ref{LB}, we present the statistical approach of Lynden-Bell \\cite{lb} to the problem of violent relaxation. In section \\ref{FD}, we show that the Fermi-Dirac equilibrium distribution predicted by Lynden-Bell is not entirely satisfactory since it has an infinite mass. We must therefore invoke {\\it incomplete relaxation} and introduce truncated models. In the artificial situation in which the system is enclosed within a spherical box, we can calculate the Fermi-Dirac spheres explicitly and prove the existence of a global entropy maximum for each value of energy. For low energies, this equilibrium state has a degenerate core surrounded by a dilute atmosphere, as calculated by Chavanis \\& Sommeria \\cite{cs}. More generally, we determine the complete equilibrium phase diagram and discuss the nature of phase transitions in self-gravitating systems. In section \\ref{MEPP}, we describe the coarse-grained relaxation of collisionless stellar systems towards statistical equilibrium in terms of a generalized Fokker-Planck equation.  This relaxation equation is derived from a phenomenological Maximum Entropy Production Principle (MEPP) and involves a diffusion in velocity space compensated by a nonlinear friction. In section \\ref{QT}, we present a quasilinear theory of the Vlasov-Poisson system and show that it leads to a kinetic equation of a Landau type.  When the system is close to equilibrium (so that a thermal bath approximation can be implemented) this equation reduces to the Fokker-Planck equation of the thermodynamical approach and the diffusion coefficient can be explicitly evaluated. This provides a new, self-consistent, equation for the ``coarse-grained'' dynamics of stellar systems where small scales have been smoothed-out in an optimal way. In section \\ref{TM}, we use this kinetic model to derive the distribution function of a tidally truncated collisionless stellar system. This truncated model preserves the main features of Lynden-Bell's distribution (including degeneracy) but has a finite mass, avoiding the artifice of a spherical container. Other truncated models attempting to take into account incomplete relaxation are discussed. ",
        "conclusions": "\\label{conclusion}   We have described in this paper the process of violent relaxation in stellar systems from the viewpoint of statistical mechanics, as originally introduced by Lynden-Bell \\cite{lb}. Lynden-Bell proposed that the structure of galaxies could result from a law of chaos: there is a total lack of information at small scales since the stars have complicated orbits, but the exciting phenomenon is that {\\it microscopic} disorder leads to {\\it macroscopic} order. The same ideas of statistical mechanics have been introduced in two-dimensional turbulence described by the Euler equation \\cite{rs1,miller,csr,csfluide} to explain the formation and maintenance of large scale coherent vortices like the Great Red Spot of Jupiter or the cyclones and anticyclones that populate the earth atmosphere \\cite{shallow,bouchet1,bouchet2}.  The formation of self-organized vortices in two-dimensional turbulence can also have applications in the context of planet formation where large-scale vortices present in the solar nebula could efficiently trap dust particles to form the planetesimals and the planets (see a complete list of references in Chavanis \\cite{cplanetes}). In fact, the statistical mechanics of the 2D Euler equation is equivalent to the theory of Lynden-Bell applying to the Vlasov equation. In a sense, the Vlasov equation is just the Euler equation for a ``fluid'' evolving in a six-dimensional phase space.  In this analogy \\cite{cthese,cfloride,japon}, the vorticity and the stream function in 2D turbulence play the same role as the distribution function and the gravitational field in galaxies. This analogy concerns not only the equilibrium states (the formation of large-scale structures) but also the relaxation towards equilibrium \\cite{csr} and the statistics of fluctuations \\cite{csire}. A kinetic theory of two-dimensional vortices can be developed in analogy with stellar dynamics \\cite{kin2}. In this kinetic theory, a point vortex experiences a diffusion process and a ``systematic drift''. This ``systematic drift'' \\cite{drift} is the counterpart of the ``dynamical friction'' \\cite{chandrabrownien} experienced by a star as a result of close encounters. Relaxation equations analogous to equations (\\ref{E25}) and (\\ref{E45}) have been derived in the context of vortex dynamics \\cite{rs,drift,kin1,kin2}. In addition, the problem of ``incomplete relaxation'' is also encountered in 2D turbulence to explain the confinement of a vortex (e.g., a dipole or a tripole) that forms after a rapid merging \\cite{csr,csfluide,rr}. It has been demonstrated explicitly in two-dimensional turbulence (where the numerical simulations are easier to implement) that the relaxation equations derived from the MEPP and including a space dependant diffusion coefficient related to the fluctuations of the vorticity (analogous to Eq. (\\ref{E65})) can account for this kinetic confinement \\cite{rr}. We believe that the relaxation equations proposed in the stellar context should work as well. The statistical mechanics of continuous vorticity fields also predicts a Fermi-Dirac distribution at equilibrium with the same interpretation of the degeneracy as in Lynden-Bell's theory. Although this degeneracy is hard to evidence for galaxies (and remains controversal), it has been vindicated by various numerical simulations and laboratory experiments of two-dimensional fluids. This suggests that the degenerate version of the theory should also be used in the stellar context.  If all these effects are taken into account correctly, it is plausible that the statistical mechanics of 2D vortices and self-gravitating systems has a chance to account for the fascinating process of self-organization in nature.       \\vskip 0.5cm"
    },
    "0212/astro-ph0212033_arXiv.txt": {
        "abstract": "A near-infrared (0.85-2.5 $\\mu$m) camera in use at the Canada-France-Hawaii Telescope and at the 1.6m telescope of the Observatoire du Mont-M\\'{e}gantic is described. The camera is based on a Hawaii-1 1024$\\times$1024 HgCdTe array detector. Its main feature is to acquire three simultaneous images at three wavelengths (simultaneous differential imaging) across the methane absorption bandhead at $1.6\\mu$m, enabling an accurate subtraction of the stellar point spread function (PSF) and the detection of faint close methanated companions. The instrument has no coronagraph and features a fast (1 MHz) data acquisition system without reset anomaly, yielding high observing efficiencies on bright stars. The performance of the instrument is described, and it is illustrated by CFHT images of the nearby star $\\upsilon$ And. TRIDENT can detect ($3\\sigma$) a methanated companion with $\\Delta H$=10 at $0.5^{\\prime \\prime}$ from the star in one hour of observing time. Non-common path aberrations between the three optical paths are the limiting factors preventing further PSF attenuation. Reference star subtraction and instrument rotation improve the detection limit by one order of magnitude. ",
        "introduction": "This paper describes a near-infrared camera built at the Laboratoire d'Astrophysique Exp\\'{e}rimentale associated with the Observatoire du Mont-M\\'{e}gantic (OMM) for use either on the 3.6m Canada-France-Hawaii telescope (CFHT) or the 1.6m telescope of the OMM. The goal was to build a compact simple instrument that would be easy to carry from the laboratory to any observatory. Commercial components were used as much as possible to minimize development time.  The camera, called TRIDENT, is based on a Hawaii-1 1024$\\times$1024 HgCdTe array detector sensitive from 0.85 to 2.5~$\\mu$m. The main scientific objective of the camera is to search for faint companions (brown dwarfs or planets) around nearby stars. Ground-based detection of faint companions is difficult because of the atmospheric turbulence that distorts the stellar diffraction pattern and optical aberrations that produce quasi-static PSF structures. TRIDENT is designed to operate with an adaptive optics (AO) system (PUEO at CFHT, \\cite{Rigaut98} or the OMM AO system\\cite{Nadeau2002,Liviu2002}) to benefit from diffraction-limited images. The camera features a special beam-splitter that allows the acquisition of three simultaneous images in three distinct narrow spectral bands across the methane absorption bandhead at $1.6\\mu$m, making possible very good subtraction of atmospheric speckles and quasi-static aberrations in the telescope and the AO system. If the three simultaneous wavelengths are carefully chosen, it is possible to enhance the star/companion contrast after image combination by selecting a spectral feature observed only in the companion but not in the star \\cite{Smith87,Rosen96}. In TRIDENT, the three wavelengths (1.567~$\\mu$m, 1.625~$\\mu$m and  1.680~$\\mu$m, 1\\% bandwidth) have been selected across the 1.6 $\\mu$m methane absorption bandhead, a spectral feature observed only in cold ($T_{\\rm{eff}}$ $<$ 1500~K \\cite{Fegley1996}) substellar objects such as T-type brown dwarfs and giant planets. TRIDENT is the practical realization of the simultaneous differential imaging concept\\cite{Racine99,Marois2000}. First light of the instrument was achieved at the Observatoire du mont M\\'{e}gantic in April 2001 (see figure 1).  \\begin{figure}[h] \\begin{center} \\includegraphics[height=13cm]{tridentfirstlight2.eps} \\caption{TRIDENT first light of Saturn at the Observatoire du Mont M\\'{e}gantic in April 2001. Upper left corner is the 1.568~$\\mu $m image, upper right is the 1.680~$\\mu $m image, lower right corner is the 1.625~$\\mu $m image and lower left is a picture of the instrument cryostat. The continuum light scattered by Saturn's ring appears equally bright in the three images while the planet is bright in the 1.568~$\\mu $m image but faint in the other two images, due to methane absorption in Saturn's atmosphere.} \\end{center} \\end{figure} ",
        "conclusions": "The TRIDENT simultaneous multi-wavelength imaging technique is a powerful way to subtract a star aberrated PSF to search for nearby faint companions. Detection thresholds are as good as or better than what has been achieved on any other telescope, showing the potential of the observing technique. Non-common path aberrations need to be kept to a minimum to assure good correlation between the three simultaneous images and achieve good subtraction performance. A new optical design is presently being integrated and will enable an improvement of the PSF subtraction by at least one order of magnitude."
    },
    "0212/astro-ph0212519_arXiv.txt": {
        "abstract": "We examine the stability of self-similar solutions for an accelerating relativistic blast wave which is generated by a point explosion in an external medium with a steep radial density profile of a power-law index $>4.134$. These accelerating solutions apply, for example, to the breakout of a gamma-ray burst outflow from the boundary of a massive star, as assumed in the popular collapsar model.  We show that short wavelength perturbations may grow but only by a modest factor $\\lesssim 10$. ",
        "introduction": "The self-similar solutions of relativistic blast waves are of much interest because of their recent applications to the study of Gamma-Ray Bursts (GRBs). A sudden release of a large amount of energy within a small volume results in a strong explosion that drives a relativistic shock into the surrounding medium. At late times the blast wave approaches a self-similar phase whereby its speed and the distribution of the pressure, density, and velocity of the gas behind the shock front do not depend on the length and time scales of the initial explosion, but only on the explosion energy and the properties of the unshocked external medium. The self-similar solutions describing this phase have been first studied by Blandford and McKee \\cite{BMK76} (hereafter, BMK).  We list their central results, which are relevant to this paper, in section \\S II.A. Note that in the BMK solution the total energy released in the explosion $E$ is the only relevant parameter.  The BMK solution is only valid for $k<4$, where $k$ is the power-law index of the radial density profile of the external medium, i.e., $\\rho_1 \\propto r^{-k}$. When $k \\ge 4$, the similarity variable defined by Blandford and McKee \\cite{BMK76} is no longer appropriate. Even in the range $3\\le k<4$, the validity of the BMK solution is not justified, because the mass contained behind the shock front diverges if the density profile of the shocked fluid is described by the BMK solution.  The self-similar solutions for steep density profiles with a power-law index $k \\ge 4$ were derived recently by Best \\& Sari \\cite{Best2000}\\footnote{The more extreme case of an exponential density profile has been discussed by Perna \\& Vietri \\cite{PV02}.}. The derivation of these solutions is similar to that in the non-relativistic regime.  The self-similar solutions to a non-relativistic blast wave were discovered independently by Sedov \\cite{Sedov46} , Von Neumann \\cite{VonNeumann47}, and Taylor \\cite{Taylor50}. The so-called \"Sedov-Von Neumann-Taylor blast wave\" solutions exist only for $k<5$, but Waxman \\& Shvarts \\cite{Waxman93} showed that in the range $3\\le k<5$ these solutions fail to describe the asymptotic flow because the energy diverges; instead they found second-type self-similar solutions for $3<k<5$ as well as for $k \\ge 5$. The new class of non-relativistic, self-similar solutions describe the flow in a limited spatial region $D(t) \\le r \\le R(t)$, where $R(t)$ is the shock radius and $D(t)$ coincides with a $C_+$ characteristic so that the flow inside the region $r< D(t)$ does not affect the flow in the outer self-similar region. The self-similar solution has to cross the sonic line into the region where the $C_+$ characteristic can not catch-up with the shock front. The solution describes a shock accelerating with a temporal dependence whose power-law index is uniquely determined by requiring that the self-similar solution cross the sonic line at a singular point. Note that in these second-type self-similar solutions the total energy released in the explosion $E$ is not a relevant parameter.  Although the energy in the self-similar part of the flow approaches a constant as time diverges, the fraction of the explosion energy $E$ carried by the self-similar component depends on the details of the initial conditions. Thus, contrary to the BMK case, dimensional arguments can not be used to determine the power-law index of the temporal dependence. Instead, the singular point determines the temporal power-law index.  In the ultra-relativistic regime, the second-type self-similar solutions for $k \\ge 4$ can similarly be obtained by requiring that they cross the sonic line at a singular point. Best \\& Sari \\cite{Best2000} found that these self-similar solutions exist for $k>5-\\sqrt{3/4}$ and describe accelerating shock waves.  However, the properties of the flow in the self-similar region, such as the energy and mass contained in the region, were not discussed.  In this paper, we rederive the self-similar solutions of ultra-relativistic blast waves for $k>4$ using a different self-similar variable and discuss the properties of the flow in the self-similar regime. Our main goal is to study the {\\it stability} of these self-similar solutions.  The stability of the Waxman-Shvarts self-similar solutions in the non-relativistic regime was studied by Sari, Waxman \\& Shvarts \\cite{Sari2000}.  They found that shocks accelerating at a rate larger than a critical value and corresponding to solutions that diverge in finite time, are unstable for small and intermediate wavenumbers.  Shocks that accelerate at a rate smaller than the critical rate are stable for most wavenumbers. The acceleration rate can be quantified by the measure $\\delta= R \\ddot{R} / \\dot{R}^2$, where the dots denote time derivatives and $R(t)$ is the radius of the shock front. This measure provides the fractional change of the velocity over a characteristic time scale for evolution ($R/\\dot{R}$). Solutions that diverge in finite time have $\\delta>1$ while others have $\\delta<1$. Thus, when shocks accelerate sufficiently fast they become unstable.  In the following sections we study the stability of the self-similar solutions of ultra-relativistic blast waves for steep density profiles with a power-law index $k>4$.  The self-similar solutions are described in \\S II. We list the BMK solutions for $k<4$ in \\S II.A and derive the self-similar solutions for $k>4$ in \\S II.B.  In \\S III, we discuss the properties of the self-similar flow and calculate the energy and mass contained in the self-similar regime. In \\S IV and \\S V, we study the stability of the self-similar solutions. Finally, we summarize our main results in \\S VI. ",
        "conclusions": "We have derived the self-similar solutions for an ultra-relativistic blast wave in an external medium with a density profile $\\rho_1 \\propto r^{-k}$ and $k>4$. The solutions exist for $k$ larger than a critical value $k_c=4.134$. They describe the flow in the self-similar region bounded by the shock front and a $C_+$ characteristic. The shock front accelerates with Lorentz factor $\\Gamma^2 \\propto t^{-m_1}$ and $m_1<-1.134$, while the $C_+$ characteristic accelerates with Lorentz factor $\\gamma^2 \\propto t^{2}$. The energy and mass contained inside the self-similar region approach constant values as time diverges.   We have found that at large wavenumbers the perturbations first decay, then grow slowly over time and eventually saturate. The initial decay and the intermediate growth are accompanied by temporal oscillations. These small wavelength perturbations grow when $1\\lesssim l/\\Gamma_0\\lesssim 10$ with an overall factor of $\\sim 10$.  At intermediate wavenumbers, the perturbations first grow slowly and then saturate. The initial growth is also accompanied by temporal oscillations. At small wavenumbers the perturbations grow monotonically in time but soon saturate. Our results also apply to expanding relativistic jets as long as the opening angle of the jet is larger than the inverse of its Lorentz factor.  In the collapsar model of gamma-ray bursts, a collimated relativistic outflow is generated due to the collapse of the core of a massive star. The outflow approaches the stellar envelope at a modest semi-relativistic speed but is expected to accelerate significantly across the sharp density gradient at the surface of the star\\cite{Zhang2002}.  Our results indicate that in the breakout phase perturbations are close to being stable in spherical symmetry.  It is still possible, however, that the lateral expansion of the jet at breakout would be accompanied by instabilities. These instabilities may produce variations in the Lorentz factor of the jet needed in the internal shock model. They may also be responsible for the complex light curves observed in most GRBs.  Current numerical simulations\\cite{Zhang2002} lack adequate resolution at the stellar surface to follow the shock breakout and confirm the instabilities. We leave a detailed study of the instabilities associated with the lateral expansion of the jet for future work.   \\paragraph*{Acknowledgments.} This work was supported in part by grants from the Israel-US BSF (BSF-9800343) and NSF (AST-0071019, AST-0204514)."
    },
    "0212/astro-ph0212315_arXiv.txt": {
        "abstract": "Observations with the Extreme UltraViolet Explorer (EUVE) satellite have shown that the inner region of the Virgo cluster (centered in M87 galaxy) has a strong Extreme UltraViolet (EUV) emission (up to $13^{\\prime}$) in excess to the low-energy tail expected from the hot, diffuse IntraCluster Medium (ICM). Detailed observations of large scale radio emission and upper limits for hard, non-thermal X-ray emission in the $2 - 10$ keV energy band have been also reported. Here we show that all available observations can be accounted for by the existence of two electron Populations (indicated as I and II) in the M87 Galaxy. The mildly relativistic Population I is responsible for the EUV excess emission via IC scattering of CBR and starlight photons. Population II electrons (with higher energy) are instead responsible for the radio emission through synchrotron mechanism. The same electrons also give rise to hard non-thermal X-ray emission (via IC scattering of CBR photons), but the resulting power is always below the upper bounds placed by present observations. The non-negligible energy budget of the two electron populations with respect to that associated with thermal electrons indicates that the M87 galaxy is not today in a quiescent (relaxed) phase. Nuclear activity and merging processes could have made available this energy budget that today is released in the form of relativistic electrons. ",
        "introduction": "Observations with the EUVE satellite (Bowyer \\& Malina \\cite{euve}) have provided evidence that a number of clusters of galaxies produce intense EUV emission, substantially in excess of the low-energy tail expected from the X-ray emission by the hot, diffuse ICM at temperature of a few keV. EUV excesses have been reported for Virgo (Lieu et al. \\cite{virgo}), Coma (Lieu et al. \\cite{coma}), A1795 (Mittaz, Lieu \\& Lockman \\cite{a1795}), A2199 galaxy clusters (Lieu, Bonamente \\& Mittaz \\cite{a2199}), A4038 (Bowyer, Lieu \\& Mittaz \\cite{blm98}) and A4059 (Durret et al. \\cite{dsldb}). However, although there is no doubt about the detection of the EUV excess emission for Virgo and Coma clusters, considerable controversy  still exists for the other clusters since EUVE results are affected by the variation of the telescope sensitivity over the field of view and depend also on subtraction of background signals that do not come from the cluster (see Bergh\\\"{o}fer \\& Bowyer \\cite{bb2002} and Durret et al. \\cite{dsldb}).  The initial explanation for the EUV excess emission was that it is produced by a warm ($T \\sim 10^6$ K) gas component of the ICM (Lieu et al. \\cite{coma}). For Virgo and Coma clusters, Bonamente, Lieu \\& Mittaz (\\cite{blm2001}) have proposed that HI cold clouds might have, at least in part, been released in the intergalactic space by ram pressure stripping. In this framework, warm gas can be generated at the interface between the cold phase (HI clouds) and the hot ICM, e.g. by the {\\it mixing layer} mechanism (Fabian \\cite{f97}). The resulting scenario is then a three phase gas model, where the hot and warm gas components are responsible for the bulk of X-ray emission and the soft excess emission, respectively, while the cold component accounts for X-ray absorption (Buote \\cite{b2001}).  Models invoking the presence in galaxy clusters of the warm gas component have problems since they refer to unrelaxed or turbulent gas conditions: the longevity of such multiphase gas would then require the presence of a heat source to carefully balance the radiative cooling, but no obvious heating mechanisms have been found (Fabian \\cite{f96}). Indeed, it is difficult for a warm intracluster gas to remain for a long time at $\\sim 10^6$K since at this temperature the cooling function for a gas with solar chemical composition has a value $\\Lambda\\simeq 2\\times 10^{-22}$ erg cm$^3$ s$^{-1}$ (Saito \\& Shigeyama \\cite{ss}) implying for the M87 galaxy warm gas component with metallicity $Z\\simeq 0.2-0.3 ~Z_{\\odot}$ (Gastaldello \\& Molendi \\cite{molendi}) and mean number density $n\\simeq 10^{-3}$ cm$^{-3}$ (Bonamente, Lieu \\& Mittaz \\cite{blm2001}) a cooling time of about 150 Myr.  Although strong evidence for cooling flows has been found in low-resolution X-ray imaging and spectra, high-resolution observations with the Reflection Grating Spectrometer on the XMM-Newton satellite show inconsistencies with the simple cooling flow models. The main problem is the lack of the emission lines expected from gas cooling below $1-2$ keV \\footnote{Several solutions to this problem have been investigated, such as heating, mixing, differential absorption and inhomogeneous metallicity (Fabianet al. \\cite{fmnp2001}), but each of these hypotheses still remain under debate (Molendi \\& Pizzolato \\cite{mp2001}). In particular, in the case of the Virgo cluster, a multiphase gas model gives a better description of the available data than a single temperature model only if it is not characterized by a broad temperature distribution, but by a narrow one (Molendi \\& Pizzolato \\cite{mp2001}). In the Virgo core (up to a radius $\\sim 10^{\\prime}$), the gas temperature ranges between 1 keV and 2 keV, so that the gas is not cold enough to justify the observed EUV emission.}. Indeed, observations with both the Hopkins Ultraviolet Telescope (Dixon, Hurwitz \\& Ferguson \\cite{dhf96}) and FUSE (Dixonet al. \\cite{dixon2001a,dixon2001b}) found no significant far-UV line emission from gas at $10^6$ K. \\footnote{For the M87 galaxy Dixon, Hurwitz \\& Ferguson (\\cite{dhf96}) found upper limits to far-UV dereddened line emission of $4.84\\times 10^{-6}$ erg cm$^{-2}$ s$^{-1}$ sr$^{-1}$ for the $O_{VI}$ line and $3.64\\times 10^{-6}$ erg cm$^{-2}$ s$^{-1}$ sr$^{-1}$ for the $C_{IV}$ line. The absence of a significant far-UV line emission, which is expected from a warm gas component, is also confirmed by FUSE observations not only for Virgo but also for the Coma cluster (Dixon et al. \\cite{dixon2001b}).} Also the observation of a large number of clusters by the XMM-Newton satellite (Tamura et al. \\cite{tamura2001,0209332}) detected no warm gas component. All these pieces of evidence seem overwhelming in the sense that a thermal mechanism for the EUV excess can be ruled out.  Therefore, other mechanisms must be investigated as the source of the EUV excess emission in clusters of galaxies. Indeed, Inverse Compton (IC) scattering of Cosmic Background Radiation (CBR) photons by relativistic electrons present in the ICM was proposed to explain the EUV excess in the case of the Coma Cluster (Hwang \\cite{hwang}, Ensslin \\& Biermann \\cite{biermann98}, Ensslin, Lieu \\& Biermann \\cite{elb99}).  Here, we will examine in detail the case of Virgo cluster, referring in particular to the results of a re-analysis of archival EUVE data for the central region of the cluster performed by Bergh\\\"{o}fer, Bowyer \\& Korpela (\\cite{bbk2000}).  In Fig. 2 of the above cited paper, the azimuthally averaged radial intensity profile of the background-subtracted EUV emission in the central region of the Virgo cluster (centered in the M87 galaxy) is given, as well as the contribution of the low-energy tail of the X-ray emitting hot ICM. Comparison between the two curves shows the presence of a diffuse EUV excess which extends up to a radius of $\\sim 13 ^{\\prime}$ (corresponding to $\\sim 65$ kpc, for an assumed M87 distance $ D = 17$ Mpc), in the M87 halo region \\footnote{ The initial analysis of the same data claimed to have detected EUV excess emission up to a distance of $20^{\\prime}$ (Lieu et al. \\cite{virgo}).}.  The aim of the present paper is to show that IC scattering of CBR and starlight photons by a population (hereafter denoted as Population I) of mildly relativistic electrons with energy up to a few hundred MeV is able to account for the EUV excess emission observed from the central region of the Virgo cluster, centered in M87 galaxy. As we will see in Section 2, Population I electrons have to be described by a power law energy spectrum with an energy cut-off at $E_{c} \\sim 250$ MeV (or slightly above) in order to avoid both a hard, non thermal, X-ray excess (due to IC scattering of CBR photons) with respect to the upper limit of $\\simeq 4\\times 10^{-12}$ erg cm$^{-2}$ s$^{-1}$ in the 2-10 keV energy band, set by observations (Reynolds et al. \\cite{rhfb}) and over-production of radio waves by synchrotron emission (Owen, Eilek \\& Kassim \\cite{owen}).  The existence in the M87 galaxy of a highly relativistic electron population (which we refer to as Population II) is well demonstrated by the observations of large-scale radio emission in the 10 MHz - 100 GHz band (Owen, Eilek \\& Kassim \\cite{owen}). This radio emission is generally interpreted as synchrotron radiation produced by relativistic electrons interacting with the halo magnetic field. To b exact, taking a value $B \\simeq 2.5~ \\mu$G as typical for the M87 halo magnetic field, one gets an electron energy range $\\sim 1 - 170$ GeV. Population II electrons have to be described with a power law energy spectrum with slope $\\alpha_e = 2.68$ in order to fit the observed radio power law spectrum $\\alpha_R = 0.84$. Electrons belonging to the same Population II, but with energy in the range $0.7 - 1.7$ GeV, also could give rise to a hard, non thermal X-ray emission in the 2 - 10 keV region (through IC scattering of CBR photons), but the resulting power is always negligible with respect to the upper limit quoted by Reynolds et al. (\\cite{rhfb}).  Regarding the origin of Population I and II electrons, it is well known that the main source of non-thermal energy in the M87 galaxy is the active nucleus and jet. The current bolometric luminosity of the inner active region is $\\sim 10^{42}$ erg s$^{-1}$ - mainly emitted in the radio band (Owen, Eilek \\& Kassim \\cite{owen}) - although the kinetic power in the jet (which matches with the nuclear power for accretion at the Bondi rate) is $\\sim 10^{44}$ erg s$^{-1}$ (Di Matteo et al. \\cite{dimatteo}). This is believed to be a lower limit for the total power available for particle acceleration, since the central engine would have had a cyclic activity with time scale of the order of 100-200 Myr and today be in a quiescent phase (Owen, Eilek \\& Kassim \\cite{owen}; Corbin, O'Neil \\& Rieke \\cite{cor}). Therefore, simple lifetime considerations indicate that the available current energy content is at least $\\sim 10^{60}$ erg. On the other hand, in the past, the M87 galaxy could have undergone a merging which would have made available a further energy budget (as suggested for the Coma cluster by Blasi \\cite{blasi}). This energy budget could be stored by accelerated protons that, being long-lived in the intergalactic medium (with lifetimes of the order of a Hubble time) could be confined in the galaxy and then released part of their energy to electrons by Coulomb and/or hadronic interactions (see e.g. Ensslin \\cite{ensslin2002}) or by shock waves in plasma turbulence.  The observed radio luminosity requires that the power emitted by Population II electrons is $\\sim 10^{42}$ erg s$^{-1}$. Since the cooling time of these electrons is dominated by synchrotron losses with a time scale of $\\sim 10^{15}$ s, one gets a Population II energy budget of $\\sim 10^{57}$ erg, much less than the current available energy content. This implies that most of the total energy made available by the M87 nucleus could be stored in Population I electrons. If these electrons are responsible for the EUV excess emission by IC scattering of CBR and starlight photons, their power is about $10^{44}$ erg s$^{-1}$, which multiplied by the IC cooling time $\\sim 10^{17}$ s, gives a Population I energy content close to that available.  As far as the electron acceleration mechanism is concerned, it is expected that Population II electrons have been directly accelerated in the active nucleus and jet with high Lorentz factors. Indeed, the two lobes in the observed radio map trace on a large-scale the jet injection, keeping initially the direction of collimated outflows from the core. Further diffusion by magnetic field inhomogeneities may give rise to a diffuse Population II component responsible for the overall radio emission. It is also known that the active nucleus and jet in the M87 inner region are responsible for inflating bubbles of accelerated cosmic rays which rise through the cooling gas at roughly half the sound speed (Churazov et al. \\cite{cbkbf}). Shock acceleration at the bubble boundaries could give rise to Population I electrons.  Finally we note that the existence for Population I electrons of an energy cutoff at $E_c \\simeq $ 250 MeV may be justified by considering that a time evolution for power-law spectrum injected electrons is expected as a consequence of transport, diffusion and cooling effects (see Petrosian \\cite{petrosian} for a detailed analysis for the Coma cluster).  In Section 2  we calculate the EUV source function for the IC scattering by relativistic electrons on CBR and starlight photons and the expected EUV excess radial profile towards the M87 galaxy. In Section 3 we present our model results and address some conclusions.    \\section {Inverse Compton Scattering}  As stated in the Introduction, we assume that Population I electrons account for the large-scale EUV excess in M87 galaxy through IC scattering of CBR and starlight photons. As far as the electron spatial distribution is concerned, despite evidence of granularity in the EUV excess map of M87 (Fig. 5 in  Bergh\\\"{o}fer, Bowyer \\& Korpela \\cite{bbk2000}), here we adopt for simplicity a continuous radial profile as representative of an azimuthally averaged radial distribution of the observed structures. Therefore, the electron density distribution $n_e(E_e,r)$ as a function of the electron energy $E_{e}$ and the radial coordinate $r$, in units of ${\\rm cm^{-3}~ eV^{-1}}$, is \\begin{equation} n_e(E_e,r) = \\displaystyle{ \\frac{K_I} {\\left[1+\\left(\\displaystyle{\\frac{r}{r_e}}\\right)^2\\right]^{3\\beta /2}} ~g(E_e)}, \\label{diste} \\end{equation} where \\begin{equation} g(E_e) = \\left\\{\\begin{array}{ll} \\left(\\displaystyle{\\frac{E_e}{{\\rm eV}}}\\right)^{-\\alpha_e} ~~~~~~~~~~~~~~~~~~~~{\\rm for}~E_e <E_c \\\\ \\\\ \\left(\\displaystyle{\\frac{E_e}{{\\rm eV}}}\\right)^{-\\alpha_e}~e^{-(E_e-E_c)/E_s} ~~~{\\rm for}~E_e > E_c~, \\end{array}\\right. \\label{gE} \\end{equation} so that the integral number density in the energy band $E_1-E_2$ can be calculated as \\begin{equation} N_e(r) = \\displaystyle{ \\int_{E_1}^{E_2} n_e(E_e,r)~dE_e }~. \\label{nde} \\end{equation} Free parameters in eqs. (\\ref{diste}) and (\\ref{gE}) are the constant $K_I$, the core radius $r_e$, the exponent $\\beta$ of the radial dependence, the electron spectral index $\\alpha_e$, the energy cutoff ($E_c$) and scale ($E_s$) parameters which, in order to avoid over production of hard X-rays and radio waves, turn out to be $E_c \\simeq 250$ MeV and $E_s \\simeq 125$ MeV, respectively. The remaining parameters have to be determined in order to fit the observed EUV excess radial profile in the energy band $50 - 200$ eV. To this end we make use of the well-known relation giving the IC scattered photon energy  (see e. g. Lang \\cite{lang}) \\begin{equation} E_{\\gamma} = \\frac{4}{3}~\\left( \\frac{E_e}{mc^2} \\right)^2 ~<E_{ph}> \\label{eqno:eic} \\end{equation} in terms of the average energy $<E_{ph}> \\simeq 8/3~$KT$_{ph}$ of target photons when a black body spectral distribution (at temperature $T_{ph}$) is assumed.  In our model the target photons are supplied by the CBR and Starlight (S) photons, with temperature $T_{CBR} \\simeq 2.7$ K and $T_{S} \\simeq 4400 $ K, \\footnote{ The temperature $T_S$ associated with the starlight black body component can be directly obtained from the photometry observations (see e.g. Marcum et al. \\cite{mof}) towards the M87 galaxy. From these observations we derive that the absolute magnitudes in the B and V bands are $M_B = -21.15\\pm 0.07$ and $M_V=-22.08\\pm 0.05$, respectively. Consequently, the starlight black body temperature turns out to be \\begin{equation} T_S = \\frac{7200~{\\rm K}}{C_{B-V}+0.68} \\end{equation} where $C_{B-V}=M_B-M_V =0.94\\pm0.09$ is the M87 color index. } respectively. According to eq. (\\ref{eqno:eic}), Population I electrons in the energy bands $125-250$ MeV and $3-10$ MeV are involved in order to scatter the CBR and starlight photons, respectively, into the EUV energy range $50-200$ eV.   The source function for EUV photon production through IC scattering is given by (Lang \\cite{lang}), in units of ${\\rm cm}^{-3}~ {\\rm s}^{-1}~ {\\rm eV}^{-1}$, as \\begin{equation} \\begin{array}{ll} Q_{IC}(E_{\\gamma},r) = \\displaystyle{ \\frac{c \\sigma_T}{2} ~ \\left( \\frac{mc^2}{{\\rm eV}} \\right) ^{1-\\alpha_e}  ~ \\left( \\frac{4}{3} \\frac{ <E_{ph}>}{{\\rm eV}} \\right) ^{\\frac{\\alpha_e-1}{2}}  } \\times \\\\ \\\\ ~~~~~~~~~\\displaystyle{ N_{ph}(r) \\frac{K_I} {\\left[1+\\left(\\frac{r}{r_e}\\right)^2\\right]^{3\\beta /2}} \\times \\left(\\frac{E_{\\gamma}}{{\\rm eV}}\\right)^{-\\frac{\\alpha_e+1}{2}} } ~, \\end{array} \\label{eqno:icsource} \\end{equation} where $N_{ph}(r)$ is the target photon number density profile, indicated as $N_{CBR}$ and $N_S(r)$ for CBR and starlight photons, respectively. Note that for $E_{\\gamma} < 200$ eV (corresponding to $E_e < E_c$) the source function $Q_{IC}(E_{\\gamma},r)$ shows a power law spectrum with slope $\\alpha_{EUV}=(\\alpha_e+1)/2$.  As far as the average background photon density $N_{ph}(r)$ in eq. (\\ref{eqno:icsource}) is concerned, in the case of CBR photons it is clearly constant ($N_{CBR}\\simeq 400$ cm$^{-3}$) within the M87 galaxy, whereas for starlight photons it is obtained by the relation \\begin{equation} N_{S}(r)  = \\frac {J_{S}(r)} {<E_{S}> c}~, \\label{ns} \\end{equation} in which the starlight energy flux $J_{S}(r)$, in units of ${\\rm eV~s^{-1}~cm^{-2}}$, is obtained from the integration of the local emissivity $ {\\cal L}_{S}(r^{\\prime})$ on the M87 volume \\begin{equation} J_{S}(r)  = \\int_{V} dV \\frac  {{\\cal L}_{S}(r^{\\prime})} {4 \\pi |r-r^{\\prime}|^2}~. \\end{equation} We describe the local emissivity by the law (Binney and Tremaine \\cite{bt}) \\begin{equation} {\\cal L}_{S}(r)  = \\frac { {\\cal L}_{S}(0) } { [1  + (r/r_{S}) ^ 2 ]^ {3/2} }~, \\label{distluce} \\end{equation} which implies that the M87 starlight luminosity profile at impact parameter $b$ is given by the modified Hubble profile \\begin{equation} I_{S}(b)  = \\frac { I_{S}(0) } { [ 1 + (b/r_{S})^2 ]} ~. \\label{distluce2} \\end{equation} In the previous two equations the relation ${\\cal L}_{S}(0) = I_{S}(0)/(2 r_{S})$ holds; moreover $I_{S}(0)$ is connected to the total luminosity through the relation \\begin{equation} I_ {S}(0) = \\frac { L_{S} } { \\pi r_{S}^2 \\log [1+(R_{M87}/r_{S})^2 ]}~. \\label{distluce3} \\end{equation} Therefore, the whole photon density profile $N_{S}(r)$ due to starlight is specified once the total luminosity $L_{S}$ and the optical core radius $r_{S}$ are given. According to Tenjes, Einasto \\& Haud (\\cite{teh}) and Zeilinger, M\\o ller \\& Stiavelli (\\cite{zms93}), we take $L_{S} = 1.07 \\times 10^{11}~L_{\\odot}$ and $r_S \\simeq 566$ pc, respectively. For definiteness we have taken the M87 galaxy size to be $R_{M87} \\simeq 100$ kpc, although our model results depend weakly  on the chosen value.  At this point we are ready to estimate the expected EUV excess radial profile in the M87 galaxy. By integrating the source function given in eq. (\\ref{eqno:icsource}) along the line of sight at impact parameter $b$ (neglecting internal absorption in M87), we get the surface brightness profile (in units of ${\\rm cm}^{-2}~{\\rm s}^{-1}~{\\rm eV}^{-1}$) \\begin{equation} \\Phi_{IC}(E_{\\gamma},b) = 2 \\displaystyle{\\int_0^{\\sqrt{R_{M87}^2-b^2}}} \\frac{Q_{IC}(E_{\\gamma},r)~r}{\\sqrt{r^2-b^2}}~dr ~. \\label{eqno:phi} \\end{equation} Correspondingly, the total power emitted in the energy band $E_1-E_2$ reads \\begin{equation} L_{IC} = 2 \\pi \\int_0^{R_{M87}} db~b~ \\int_{E_1}^{E_2} dE_{\\gamma}~\\Phi_{IC}(E_{\\gamma},b)~E_{\\gamma} ~. \\end{equation}  The photon number count, in units of ${\\rm arcmin^{-2}~s^{-1}}$, expected to be measured by the Deep Survey telescope  aboard the EUVE satellite as a function of the angular radius $\\theta \\simeq b/D$ results to be \\begin{equation} {\\cal N}_{EUV}(\\theta) = \\displaystyle{\\int_{E_1}^{E_2}} \\Phi_{IC}(E_{\\gamma},\\theta D) e^{-\\tau(E_{\\gamma})}~A_{e}(E_{\\gamma})dE_{\\gamma}~, \\label{eqno:counts} \\end{equation} where $[E_1-E_2] \\simeq [50-200]$ eV is the sensibility energy range of the DS instrument and $A_{e}(E_{\\gamma})$ its effective area given in Bowyer et al. (\\cite{aeff}). In eq. (\\ref{eqno:counts}) we also take into account the galactic absorption through the optical depth $\\tau(E_{\\gamma}) = N_H \\sigma(E_{\\gamma})$, $N_H = 1.7 \\times 10^{20}$ cm$^{-2}$ being the Milky Way column density towards M87 (Hartman \\& Burton \\cite{hb97}) and $\\sigma(E_{\\gamma})= [11~\\sigma_{\\gamma H}(E_{\\gamma}) + \\sigma_{\\gamma He}(E_{\\gamma})] /12$ is accounting for both the cross section of $\\gamma H$ and $\\gamma He$ interaction, as parameterized by Morrison \\& McCammon (\\cite{mmc83}) and Yan, Sadeghpour \\& Dalgarno (\\cite{ysd98}), respectively. ",
        "conclusions": "Following the formalism discussed in Section 2, we fit the available EUV excess data towards the M87 galaxy. For this purpose, experimental points (given with the respective error bars in Fig. \\ref{fit}) are derived from those given in Fig. 2 in Bergh\\\"{o}fer, Bowyer \\& Korpela (\\cite{bbk2000}), by subtracting the contribution of the low-energy tail due to the X-ray emitting ICM.  Preliminarily, we have fitted the M87 observed EUV excess by considering only the IC scattering of Population I electrons on CBR photons  (Model I, with 11 d.o.f.). In this way, we obtain a poor fit with $\\chi^2=2.3$ by considering all the 15 observed points (see the dotted line in Fig. \\ref{fit}). \\footnote{The $\\chi^2$ value gets reduced to $\\chi ^2 \\simeq 0.7$ if we do not include the first observed point at $1^{\\prime}$. Indeed, a substantial fraction of the observed number count at $1^{\\prime}$ may be the result of the core and jet activity (Bergh\\\"{o}fer, Bowyer \\& Korpela \\cite{bbk2000}).}  Consideration of the IC scattering by the same electron population also on starlight photons, distributed accordingly to eqs. (\\ref{ns}) - (\\ref{distluce3}),  (Model II, with 9 d.o.f.), allows us to substantially increase the expected photon number counts from the innermost part of the galaxy, therefore obtaining a much better confidence level $\\chi ^2 \\simeq 0.5$, considering all the 15 points (see solid line in Fig. \\ref{fit}). This represents a significant improvement of the fit to all the observed points since the outcome of the F-test implies that Model II has to be preferred over Model I within a $3\\%$ confidence level (the F-test value is 3.9).  Selected models with $\\chi^2 < 1$ give a range of values for $r_e$ and $\\beta$ given by $r_e = 5 \\pm 1$ kpc and $ \\beta = 0.5 \\pm 0.1$. As far as the remaining two parameters $\\alpha_e$ and $K_I$ - which are actually related by eq. (\\ref{nde}) - our acceptable $\\chi^2 < 1$ fits give $ 2 < \\alpha_e < 3$ and, correspondingly, $10^4  \\ut < K_I  \\ut <  10^8$ eV$^{-1}$ cm$^{-3}$. The obtained relative ratio between the EUV luminosity from IC scattering on starlight and CBR photons is in the range of $2\\%-6\\%$. The parameter $\\alpha_e$ cannot be further constrained by EUV excess data due to the absence of spectral measurements.  Recent experimental results (Reynolds et al. \\cite{rhfb}) set an upper limit of $4 \\times 10^{-12}$ erg s$^{-1}$ cm$^{-2}$ for the hard, non-thermal X-ray flux in the 2-10 keV energy band from the M87 galaxy. Moreover, radio observations in the 10 MHz - 100 GHz band give a total luminosity $L_R \\simeq 5 \\times 10^{41}$ erg s$^{-1}$ (Herbig \\& Readhead \\cite{herbig}, Owen, Eilek \\& Kassim \\cite{owen}). If we do not introduce an energy cutoff for Population I electrons (see eq. (\\ref{gE})), we would find both a hard, non-thermal X-ray excess emission (in the 2-10 keV energy band, due to IC scattering of CBR photons) and radio wave excess by synchrotron mechanism, above the upper bound allowed by observations.  In the following, as stated in the Introduction, we assume that the radio emission is accounted for by Population II electrons described by the same distribution law as in eq. (\\ref{diste}) for Population I electrons, but with different parameter values for $K_{II}$, $r_e$, $\\beta$, $\\alpha_e$ and without energy cutoff. \\footnote{Indeed, Population I and II electrons are expected to have different spatial distributions although the radio map of the M87 halo shows some similar features with respect to the observed distribution of the outer EUV excess (see Bergh\\\"{o}fer, Bowyer \\& Korpela (\\cite{bbk2000}), Andernach, Baker, van Kap-herr \\& Wielebinski (\\cite{abkw}) and Churazov et al. (\\cite{cbkbf})).} Here we recall that the characteristic frequency and the differential luminosity of the ultrarelativistic synchrotron radiation are given by (see e. g. Lang \\cite{lang}) \\begin{equation} \\nu = \\frac{3}{4 \\pi}\\frac{eB}{mc}\\left( \\frac{E_e}{mc^2}\\right)^2~ \\label{eqno:syn} \\end{equation} and \\begin{equation} L_R(\\nu) =  \\frac{ \\sqrt{3} e^3 C_R F_e }{8 \\pi m c^2} \\left( \\frac{3 e}{4 \\pi m^3 c^5}\\right) ^{\\frac{\\alpha_e-1}{2}} B^{\\frac{\\alpha_e+1}{2}} ~ \\nu^{-\\frac{\\alpha_e-1}{2}}~, \\label{eqno:nusyn} \\end{equation} where $C_R \\simeq (1.602 \\times 10^{-12}~{\\rm erg})^{\\alpha_e}$ and \\begin{equation} F_e =  \\displaystyle{   \\int_0^{R_{M87}} \\frac{K_{II}} {\\left[1+\\left(\\frac{r}{r_e}\\right)^2\\right]^{3\\beta /2}} ~4 \\pi r^2~dr}~. \\label{fe} \\end{equation} If we take $\\alpha_e \\simeq 2.68$, so that the predicted radio power law spectrum $\\alpha_R = (\\alpha_e - 1)/2 \\simeq 0.84$ is in agreement with the observed value, and a value $B \\simeq 2.5~\\mu$G for the magnetic field strength in the M87 halo (Dennison \\cite{d80}), there exists a set of acceptable values for the free parameters in eq. (\\ref{fe}) for which the obtained diffuse radio emission from the M87 galaxy in the band 10 MHz -- 100 GHz is $L_{R} \\simeq 5 \\times 10^{41}$ erg s$^{-1}$, consistent with the observed value (Herbig \\& Readhead \\cite{herbig}, Owen, Eilek \\& Kassim \\cite{owen}).  It is remarkable that the same set of values for Population II electron parameters that fit radio observations give a hard, non-thermal X-ray flux by IC scattering of CBR photons (by electrons in the energy range 0.7 - 1. 7 GeV) in accordance with the upper bounds given by Reynolds et al. (\\cite{rhfb}) for the M87 galaxy. \\footnote{Of course, this estimate is obtained by using the same formalism described in Section 2 for the EUV excess by Population I electrons, but using in eq. (\\ref{eqno:icsource}) the appropriate Population II electron parameter values.} More recently, the XMM-Newton observatory has been pointed towards M87 galaxy with the aim of investigating if non-thermal IC emission is present in the arms regions and if it is linked to the power emitted by the radio jet. The upper limit to the non-thermal X-ray emission has been found to be $\\leq 4\\times 10^{-14}$ erg cm$^{-2}$ s$^{-1}$ in the $0.5-8$ keV band, representing less than $1\\%$ of the flux from the thermal components (Belsole et al. \\cite{belsole}). Our model results give a hard non-thermal X-ray flux from the inner $\\simeq 5\\arcmin$ region of M87 galaxy close to the above upper bound. Clearly, our model results depend on the chosen parameters for population II electrons which are not, at present, tightly determined by radio observations. For example, by increasing the Population II electron core radius (or decreasing the value of $\\beta$) we can dilute the X-ray excess (and also the radio emission) within a wider region, leaving the same total power.  \\begin{figure}[htbp] \\vspace{7.7cm} \\special{psfile=ms2035f1.eps vscale=38. hscale=46. voffset=-50 hoffset=-20} \\caption{ The expected photon number count ${\\cal N}_{EUV}$ is shown as a function of the angular radius $\\theta$, for the model giving the best-fit (solid line) to the observational data. Dotted and dashed lines represent the contribution from IC scattering of CBR and starlight photons, respectively. } \\label{fit} \\end{figure}  A few comments about the energy budget involved for the two electron Populations are in order. Observations show that the current bolometric luminosity of the M87 inner active region is $\\sim 10^{42}$ erg s$^{-1}$ - mainly emitted in the radio band (Owen, Eilek \\& Kassim \\cite{owen}) - although the kinetic power in the jet is $\\sim 10^{44}$ erg s$^{-1}$ (Di Matteo et al. \\cite{dimatteo}). Simple lifetime considerations then indicate that the available electron energy content is at least $\\sim 10^{60}$ erg. However, the total power available for particle acceleration may be larger since the central engine could have had a cyclic activity with a time scale of the order of 100-200 Myr and today be in a quiescent phase (Owen, Eilek \\& Kassim \\cite{owen}; Corbin, O'Neil \\& Rieke \\cite{cor}). Moreover, in the past the M87 galaxy could have undergone a merging phenomenon which would have made available a further energy budget, initially stored by relativistic protons and later released (partially) to electrons by Coulomb and/or hadronic interactions or by shock waves in plasma turbulence (Ensslin \\cite{ensslin2002}).  Coming back again to our model results, we find that the energy budget associated with the two electron Populations is about $10^{60}$ erg (mostly in Population I electrons), which is always less or approximately equal to the total energy associated with thermal electrons. \\footnote{The energy associated with thermal electrons is estimated by assuming that the hot, virialized, X-ray emitting gas has a solar metallicity composition and is distributed following a $\\beta$-model with $\\beta \\simeq 0.5$, core radius $\\simeq 1.6$ kpc and temperature $T_g \\simeq 2.5$ keV (Fabricant and Gorenstein \\cite{fg}). Moreover, the gas central number density is taken to be $\\sim 3 \\times 10^{-1}$ cm$^{-3}$ so that the total X-ray luminosity in the energy bands 0.2 - 4 keV and 2 - 10 keV turns out to be $\\simeq 2.7 \\times 10^{43}$ erg s$^{-1}$ and $\\simeq 3 \\times 10^{43}$ erg s$^{-1}$, in accordance with the observed values (Forman and Jones \\cite{fj}).} This result represents an indication that the M87 galaxy is not today in a quiescent (relaxed) phase.  However, further more accurate spectroscopic observations in both the EUV and X-ray bands (especially towards the external region of the M87 galaxy) are necessary in order to confirm the presence of the two electron Populations invoked in the present paper. Moreover, observations towards the other galaxy clusters for which the EUV emission is still uncertain may establish whether the existence of these two electron Populations is a common feature in galaxy clusters or if it is particular to the central, massive, cD galaxies like M87."
    },
    "0212/astro-ph0212123_arXiv.txt": {
        "abstract": "Using the {\\it Echellete Spectrograph Imager} (\\esi) on the Keck~II 10-m telescope we have measured the redshifts of the host galaxies of gamma-ray bursts GRB\\,990506 and GRB\\,000418, $z=1.30658 \\pm 0.00004$ and $1.1181 \\pm 0.0001$, respectively.  Thanks to the excellent spectral resolution of \\esi\\ ($\\lambda/\\Delta\\lambda = 13000$) we resolved the \\hbox{[O II]} 3727 doublet in both cases. The measured redshift of GRB\\,990506 is the highest known for a dark burst GRB, though entirely consistent with the notion that dark and non-dark bursts have a common progenitor origin. The relative strengths of the \\hbox{[O II]}, \\hbox{He I}, \\hbox{[Ne III]}, and \\hbox{H$\\gamma$} emission lines suggest that the host of GRB\\,000418 is a starburst galaxy, rather than a LINER or Seyfert 2.  Since the host of GRB\\,000418 has been detected at sub-millimeter wavelengths these spectroscopic observations suggest that the sub-millimeter emission is due to star-formation (as opposed to AGN) activity.  The \\hbox{[O II]}-derived unobscured star-formation rates are 13 and 55 $M_\\odot$ yr$^{-1}$ for the hosts of GRB\\,990506 and GRB\\,000418, respectively. In contrast, the star-formation rate of the host of GRB\\,000418 derived from sub-millimeter observations is twenty times larger. ",
        "introduction": "The determination of redshifts continue to play a crucial r\\^{o}le in the understanding of the physics of GRBs.  Without redshifts, we are unable to translate observed fluences and angular burst-host offsets into physically meaningful quantities.  That redshift is an integral part of the physical modeling of the afterglow is best demonstrated by the immense difficulty in modeling the afterglow of GRB\\,980329 \\citep{yfh+02}, a GRB that enjoys an abundance of pan-chromatic afterglow observations but unfortunately lacked a measurement of redshift.  Emission spectra of GRB hosts themselves are valuable given the potential use of GRBs to infer the star-formation history of the Universe. Since $\\gamma$-rays penetrate dust, a GRB-defined sample does not suffer from the well known biases of optical/UV samples. In addition, the high-accuracy astrometric localizations afforded by radio observations (with, e.g., the Very Large Array; VLA) and/or X-ray (Chandra) offer significant observational advantages over a sub-millimeter-defined sample. Indeed, one great hope for the utility of GRBs is in revealing the star-formation history at extremely high redshifts \\citepeg{bl02} since the brilliance of the prompt emission allows these events to be seen to redshifts well beyond 10 \\citep{lr00}.  GRB\\,990506, is a canonical ``dark burst,''\\footnotemark\\footnotetext{GRB\\,970828 is another such dark burst, with a similar detection history as GRB\\,990506 \\citep{dfk+01}: later optical observations showed that the radio-localized host of GRB\\,970828 was a dusty galaxy at $z\\sim 1$.} with a radio-detected afterglow with no optical transient emission \\citep{tbf+00}.  Since no transient optical emission was detected from the afterglow, an optical absorption-line redshift of GRB\\,990506 was impossible to obtain.  \\citet{tbf+00} discovered an apparent two-galaxy system in optical data from Keck consistent with the radio transient position.  Subsequent imaging from the {\\it Hubble Space Telescope} (HST) revealed the putative two-galaxy system to be morphologically distinct \\citep{htab+00}. \\citet{bkd02} later refined the radio$\\rightarrow$HST astrometry to show that the radio afterglow position is only consistent with the southwest component; this component is identified as the host of GRB\\,990506. However, given the poor astrometric tie between the radio and the optical positions, as noted by \\citeauthor{bkd02}, the GRB\\,990506--galaxy association is one of the more tenuous identifications of a host galaxy.  GRB\\,000418 is famous for two reasons. First, it is still one of the best cases for collimated ejecta of a GRB explosion \\citep{bdf+01}. More important, the host galaxy has been detected in the radio (using the Very Large Array, VLA) and the sub-millimeter bands (using the Sub-millimeter Common User Bolometer Array, SCUBA, on the James Clerk Maxwell Telescope); see \\citet{eck+02}.  If the sub-mm emission arises from reprocessing of stellar UV light by dust (as opposed to arising from AGN activity) then the inferred star-formation rate is very high. If so, this galaxy can potentially provide information about the nature of these ultra-luminous galaxies and their contribution to the global star formation rate.  Here we present the spectroscopic observations of the hosts of GRB\\,000418 and GRB\\,990506. The discussion of GRB\\,000418 here expands our preliminary announcement of the redshift of the host \\citep{bdd+00}.  In \\S \\ref{sec:esiobs} we describe the observations and reductions that led to the redshift determinations (presented in \\S\\S \\ref{sec:0506redshift}--\\ref{sec:0418redshift}).  In \\S \\ref{sec:natureofhosts} we discuss some of the implications of the redshift detections in the context of GRB host galaxies. ",
        "conclusions": "Traditionally, low-resolution spectroscopy ($R \\approx 1000$) has been used to determine emission-line redshifts of GRB hosts. Indeed, on the night previous to the [O II] detection of the host of GRB\\,990506 in echellete mode, a one hour low-resolution long-slit prism spectrum was obtained and no lines were detected. This non-detection motivated our subsequent use of the spectrograph in echellete (medium resolution) mode.  Medium- to high-resolution spectrometers have already been used for the detection of absorption-line redshifts of GRBs \\citep{cdd+00,cdk+00,mhk+02,sfi+02}, but the redshifts of the hosts of GRB\\,000418 and GRB\\,990506 are the first emission-line redshifts found with a medium-resolution ($R=13000$) echellete spectrograph.  This point is worth highlighting from an observational perspective. Given that the throughput of \\esi\\, in echellete mode ($\\ale 12\\%$) is approximately 40\\% smaller than that of the prism long-slit (``high-throughput'') mode \\citepcf{sbe+02}, it is not immediately obvious why echellete spectroscopy would be able to detect the redshift of GRB\\,990506 whereas long-slit spectroscopy failed to detected the line.  While both observing modes provide similar wavelength coverage and slit widths, however, the prism long-slit mode results in a lower dispersion with increasing wavelengths (from 1 \\AA/pixel at 4000 \\AA\\ to $\\sim$10 \\AA/pixel at 10,000 \\AA).  This effect causes the numerous sky emission lines to overlap to the extent that, beyond about 6000 \\AA, almost all pixels are dominated by dispersed sky emission lines.  In the echellete-mode, the sky lines are always resolved except for the largest slit widths.  Even if sky subtraction were perfect, the noise from the high sky background levels in low-resolution mode more than mitigates, in terms of signal--to--noise, against the gain in throughput.  Is it clear from previous studies that GRB hosts contain moderate levels of unobscured star formation, resulting in strong Balmer line, [O II] and Ly $\\alpha$ emission-lines \\citep[see][for review]{dkb+01}. Most GRBs, however, appear to occur at redshifts near unity and so H$\\alpha$ and Ly$\\alpha$ lines are outside the range of detectability for optical CCDs. At such redshifts, [O II], typically the third most luminous star-formation line, resides at $\\lambda \\age 6000$ ($z \\age 0.6$) in the observer's frame; however, this is the onset of the wavelength regime where the number of prominent sky lines becomes numerous and densely packed.  As a faint, typically narrow line (observed FWHM $\\ale 6$ \\AA), the [O II] line can be easily outshone from a nearby night sky line when the dispersion is low.  This was clearly the case in GRB\\,990506 (Figure~\\ref{fig:spect0506}).  Medium-resolution spectroscopy also affords a better insight into the nature of any emission line detections.  Specifically, the ambiguity of single line redshifts \\citep[see][]{sbs+00} is essentially removed since the resolution of \\esi\\, is large enough to resolve the [O II] doublet (restframe separation of 2.75 \\AA) for redshifts greater than $z \\sim 0.5$.  From the optical lines, we have shown that the unobscured star-formation rates in the hosts of GRB\\,990506 and GRB\\,000418 are $13$ and $55$ M$_\\odot$ yr$^{-1}$, respectively.  This may be contrasted with ${\\rm SFR}\\sim 600$ M$_\\odot$ yr$^{-1}$ from sub-mm and radio observations of GRB\\,000418.  Such a discrepancy has been observed in many high-$z$ galaxies and it is probably due to dust obscuration. Such a trend, for GRB galaxies, has been previously noted by comparing infrared photometry with spectroscopic star-formation indicators \\citep{cba02}. More importantly, however, spectroscopy of GRB\\,000418 has allowed us to determine the redshift to a sub-mm galaxy.  This is a significant result since there are only a handful of such galaxies with a measured redshift \\citepeg{lso+02}.  This hints at the unique potential of GRB-selected galaxies in uncovering the redshift distribution of the population of sub-mm galaxies.  Since GRB\\,990506 is one of only a handful of well-studied dark bursts (also, GRB\\,000210, \\citealt{pfg+02} and GRB\\,970828, \\citealt{dfk+01}) the measurement of a redshift near the median of other long-duration bursts may suggest a common progenitor population. This emerging trend, of similar redshift distributions, will almost certainly be testable with the advent of systematic absorption redshift determinations for {\\it Swift} bursts."
    },
    "0212/astro-ph0212409_arXiv.txt": {
        "abstract": "An analysis of a 29 ksec {\\it Chandra} ACIS-S observations of the young, Cassiopeia A-like supernova remnant in the irregular galaxy NGC 4449 is presented. The observed $0.5 - 2.1$ keV spectrum reveals the likely presence of several emission lines including \\ion{O}{8} at 0.65 and 0.77 keV, \\ion{Ne}{10} at 1.05 keV, \\ion{Mg}{11} at 1.5 keV, and \\ion{Si}{13} at 1.85 keV. From the observed spectrum, we derive an N$_{\\rm H}$ = 1 $\\times$ $10^{21}$ cm$^{-2}$ and an X-ray temperature of T $\\approx$ 9 $\\times$ 10$^{6}$ K. A non-equilibrium ionization fit to the spectrum suggests an overabundance of oxygen around 20 times solar, consistent with the remnant's UV and optical emission-line properties. We discuss the remnant's approximate X-ray derived elemental abundances and compare its X-ray spectrum and luminosity to other oxygen-rich remnants. ",
        "introduction": "At present, there are only about a half-dozen O-rich supernova remnants (SNRs) known. They represent the  remains of high-mass stars (M $\\gtrsim$ 15 M$_{\\sun}$) and thus are especially useful for studying the elemental abundances of explosive nucleosynthesis processes in core-collapse supernovae. The younger the remnant is, the more likely its properties accurately reflect true SN ejecta abundances unaffected by dilution from swept-up circumstellar and interstellar material.  Currently, the youngest Galactic remnant is the $\\simeq$ 320 yr old O-rich SNR Cassiopeia A (Cas~A). However, with an estimated age of around 100 yr (\\citealt{blair83,blair98}), the O-rich SNR located in the Magellanic type irregular galaxy NGC 4449 represents the youngest example of this type of remnant and is therefore an object of particular interest.  The NGC 4449 SNR was first identified by \\citet{seaquist78}, who found a strong and point-like, non-thermal radio source approximately 1\\arcmin~north of the nucleus of this galaxy coincident with an \\ion{H}{2} region cataloged by \\citet{sabbadin79}. Optical spectrophotometry of this source by \\citet{balick78} showed two components: (1) narrow lines belonging to a conventional \\ion{H}{2} region, and (2) broad [\\ion{O}{3}] $\\lambda \\lambda$4959,5007 and [\\ion{O}{1}] $\\lambda \\lambda$6300,6363 line emissions which they attributed to a young SNR similar to Cas A.  Optical and UV spectra taken by \\citet{kirshner80}, \\citet{blair83}, and \\citet{blair84} confirmed and extended this interpretation and yielded a rough analysis of the physical conditions in the SNR and the associated \\ion{H}{2} region. They found optical line widths implying an expansion velocity of 3500 km s$^{-1}$ and emphasized the extremely low abundance of hydrogen in the ejecta. Furthermore, they estimated that the total mass of oxygen in the optical was 10$^{-2}$ $M_{\\sun}$, some 50 times higher than the value derived for Cas A \\citep{peimbert71}, and estimate a progenitor mass $\\approx$ 30 $M_{\\sun}$.  Using data from the {\\it Einstein X-Ray Observatory} High Resolution Imager, \\citet{blair83} found the SNR to also be highly luminous in X-rays, with L$_{\\rm x}$ $\\approx$ 10$^{39}$ erg s$^{-1}$ assuming a distance of 5.0 Mpc.  Even at a smaller estimated distance of 3.9 Mpc \\citep{devaucouleurs91,hunter99}, this SNR is nearly 10 times more luminous in X-rays than Cas A, the most luminous young SNR in our Galaxy, and several times more luminous than the bright O-rich SNR N132D in the Large Magellanic Cloud (L$_{\\rm x}$ $\\approx$ 6 $\\times$ 10$^{37}$ erg s$^{-1}$; \\citealt{hughes87}).  More recently, \\citet{vogler97} studied {\\it ROSAT} ~PSPC and HRI detected NGC~4449 sources, including the SNR. They were able to fit the remnant's spectrum to an absorbed thermal bremsstrahlung model, reporting a lower limit on both the absorbing column $N_H \\geq$ 7 $\\times$ 10$^{20}$ cm$^{-2}$ and temperature of $\\geq$ 600 eV corresponding to T $\\geq$ 7 $\\times$ 10$^{6}$ K. They also found a L$_{\\rm x}$ of $\\approx$ 1.8 $\\times$ 10$^{38}$ erg s$^{-1}$, a value about 50\\% that of the Blair et al. (1983) {\\it Einstein} 1978 estimate.  Here we report on observations of the SNR in NGC 4449 obtained with the {\\it Chandra X-ray Observatory} \\citep{weisskopf96}. In \\S \\ref{sec:obs} we describe the observations and data reduction. In \\S \\ref{sec:comp} we discuss our model fits to the data, the derivation of various physical quantities associates with the X-ray emitting plasma, and how the SNR compares to other O-rich SNR. Our results and conclusions are summarized in \\S \\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  We have analyzed a 29 ksec {\\it Chandra} ACIS-S observation of the  $\\approx$ 100 yr old O-rich supernova remnant in the irregular galaxy NGC 4449. We found the following:  1) The spectrum shows the likely presence of several emission lines including \\ion{O}{8}, \\ion{Ne}{10}, \\ion{Mg}{11}, and \\ion{Si}{13} with an overall spectrum similar in appearance to the oxygen-rich Galactic SNR G292.0+1.8.  2) We find an X-ray luminosity of 2.4 $\\times$ 10$^{38}$ erg s$^{-1}$, for the energy range 0.5 -- 2.1 keV and a distance of 3.9 Mpc. Assuming a spherical geometry for the remnant we estimate the mass of the X-ray emitting gas, including the mass of metals as well as He and H, to be M$_{\\rm Xray}$ $\\approx$ 5.7 {\\it f} D$^{2.5}_{\\rm 3.9 Mpc}$ M$_{\\sun}$.  3) A non-equilibrium ionization fit to the spectrum suggests a large overabundance of oxygen of $\\approx$ 18 times solar, consistent with the remnant's UV and optical emission-line properties.  4) Based on electron-ion temperature equilibration, we derive an X-ray shock velocity of $\\approx$ 830 km s$^{-1}$. This is significantly lower than the 6000 km s$^{-1}$ optical expansion velocity determined from optical measurements but consistent with differences between the optical and X-ray velocities seen in other O-rich remnants.  5) While our estimated  X-ray flux is in rough agreement with that measured by {\\it ROSAT} in 1994, it is only about half of that of an 1978 {\\it Einstein} observation. Despite uncertainties in the accuracy of such flux measurement comparisons, such a decrease would not be unexpected given recent measurements made in the radio and optical suggesting that the remnant continues to undergo significant changes since it discovery in 1978.  The NGC 4449 SNR's current standing as the youngest and most luminous O-rich remnant known certainly make it an object worthy of continued study. The remnant's estimated angular dimensions of $\\approx$ $0 \\farcs 028$ make it an ideal candidate for VLBI measurements of its true size and emission sub-structure, like that accomplished for SN 1993J \\citep{bietenholz01}. Such measurements could help clarify its general structure and, when combined with optical emission line width data, set strict limits on its expansion age."
    },
    "0212/astro-ph0212253_arXiv.txt": {
        "abstract": "A new measurement of the primary cosmic--ray proton and helium fluxes from 3 to 350~GeV was carried out by the balloon--borne CAPRICE experiment in 1998. This experimental setup combines different detector techniques and has excellent particle discrimination capabilities allowing clear particle identification. Our experiment has the capability to determine accurately detector selection efficiencies and systematic errors associated with them. Furthermore, it can check for the first time the energy determined by the magnet spectrometer by using the Cherenkov angle measured by the RICH detector well above 20~GeV/n. The analysis of the primary proton and helium components is described here and the results are compared with other recent measurements using other magnet spectrometers. The observed energy spectra at the top of the atmosphere can be represented by (1.27$\\pm$0.09)~$\\times$~10$^{4}$~E~$^{-2.75\\pm0.02}$~particles~(m$^2$~GeV~sr~s)$^{-1}$, where E is the kinetic energy, for protons between 20 and 350~GeV and (4.8$\\pm$0.8)~$\\times$~10$^{2}$~E~$^{-2.67\\pm0.06}$~particles~(m$^2$~GeV~nucleon$^{-1}$~sr~s)$^{-1}$, where E is the kinetic energy per nucleon, for helium nuclei between 15 and 150~GeV~nucleon$^{-1}$. ",
        "introduction": "Accurate measurements of the spectra of primary cosmic rays have deep astrophysical implications, since they provide important information on the mechanisms of production of cosmic rays and of the matter distribution in the interstellar space. In fact, the spectral shapes of proton and helium nuclei fluxes are sensitive indicators of the processes of particle acceleration, and the observed fluxes are the primary measure of the energy density of cosmic rays in the interstellar medium. Their spectra also serve as important inputs to calculations that aim to predict the secondary antiproton or positron spectra, which result from high--energy interactions of protons and helium nuclei with the interstellar gas.  Recently, the importance of the normalization of the primary cosmic--ray flux has been emphasized in connection to the atmospheric neutrino observations performed by underground experiments, e.g. Super--Kamiokande \\cite{fukuda}. A correct interpretation of these measurements depends on the accuracy of the predictions to which they are compared. The assumptions about the flux of cosmic rays that impinge on the Earth turn out to be among the main sources of inaccuracies in  the simulation of atmospheric showers \\cite{gaisser01}. For energies above 10~GeV, where the effect of the solar modulation is less than about 10\\%, there is still a difference of about 10\\%--20\\% in the absolute fluxes published in the recent years \\cite{boezio99b,sanuki,alcaraz,seo,bellotti,menn}.  From this experimental scenario, the need arises for measurements extended over a large energy range and with a good understanding of the systematic uncertainties associated with them. Measurements of primary particles have been carried out using different techniques: magnet spectrometers, e.g. \\cite{seo,bellotti}, and RICH detectors \\cite{buckley,diehl} have been used for energies up to 100--200 GeV/n, while calorimetric measurements extend to higher energies, e.g. \\cite{silverberg,asakimori}. However, comparisons among such measurements show sometimes significant discrepancies.  The differences are probably due to imprecise or incomplete knowledge of detectors response functions during the flight. Monte~Carlo simulations or calibrations at low energy in the laboratory prior to flight are not ideal since experimental conditions in the gondola such as temperature or pressure may not be stable during the flight. If the response function is based on the performance of individual detectors in the laboratory, systematic uncertainties are likely when operated with other detectors and would introduce a bias. So, the best experimental determination of the efficiencies and calibrations is to make use of redundant detectors to select a set of ``good'' events independent of the other detectors. These sets of events can be used to determine the response function of the detectors which are not involved in the selection.  In this paper we report a new measurement of the primary proton and helium nuclei spectra with the CAPRICE98 instrument, which combines different detector techniques. This instrument, described in section \\ref{c98}, consisted of a superconducting magnet spectrometer, a Ring--Imaging Cherenkov detector and an imaging calorimeter. With these independent detectors, it was possible to accurately determine both the efficiency and the systematic error associated with each detector. The data analysis is described in section \\ref{datana}, which includes the determination of detector efficiencies, systematic errors and the corrections applied. The final results and discussions are presented in section \\ref{finres}. The results cover a range of kinetic energy from 3 to 350~GeV for protons, and from 0.9 to 150~GeV~nucleon$^{-1}$ for helium nuclei. ",
        "conclusions": "The primary proton and helium nuclei results from the latest CAPRICE balloon--borne experiment, performed in 1998, were presented. The excellent performance of this apparatus allowed an accurate measurement of the spectra extended over a large energy range. For the first time it was possible to cross--check the rigidity measurement during the flight at high energy, where the measured spectra are more sensitive to this kind of systematic errors. The CAPRICE98 instrument made it possible to accurately determine efficiencies, rejection power and to estimate systematic errors for each individual detector. This allowed us to measure proton and helium nuclei spectra with an excellent understanding of the performance of the detectors. The results are in good agreement with other recent measurements if the systematic errors are properly taken into account. However, all these results are significantly lower than some of the older measurements.  New space experiments, PAMELA \\cite{pamela} in 2003 and AMS2 \\cite{ams2} later on the International Space Station, will be able to perform the same measurements with better statistics than was done in the past, but the CAPRICE98 instrument will remain a unique detector that joined together a precise superconducting magnet spectrometer and a high threshold gas RICH detector, an excellent apparatus to study\\linebreak cosmic--rays.  \\ack This work was supported by NASA Grant NADW--110, the Istituto Nazionale di Fisica Nucleare, Italy, the Agenzia Spaziale Italiana, DARA and DFG in Germany, the Swedish National Space Board and the Knut and Alice Wallenberg foundation. The Swedish--French group thanks the EC SCIENCE programme for support. P.H. was supported by the Swedish Foundation for International Cooperation in Research and Higher Education; E.M. was supported by the Foundation BLANCEFLOR Boncompagni--Ludovisi, n\\'ee Bildt. We wish to thank the National Scientific Balloon Facility and the NSBF launch crew that served in Fort Sumner. We would also like to acknowledge the essential support given by the Gas Work Group (EST/SM/SF) and the Thin Films \\& Section (EP/TA1/TF) at CERN, the LEPSI and CRN--Strasbourg and the technical staff of NMSU and of INFN.  \\newpage"
    },
    "0212/hep-ph0212146_arXiv.txt": {
        "abstract": "The first results from the KamLAND experiment have provided confirmational evidence for the Large Mixing Angle (LMA) Mikheyev-Smirnov-Wolfenstein (MSW) solution to the solar neutrino problem. We do a global analysis of solar and the recently announced KamLAND data (both rate and spectrum) and investigate its effect on the allowed region in the $\\Delta m^2-\\tan^2\\theta$ plane. The best-fit from a  combined analysis which uses the KamLAND rate plus global solar data comes at $\\Delta m^2 = 6.06 \\times 10^{-5}$ eV $^2$ and $\\tan^2\\theta=0.42 $, very close to the global solar best-fit, leaving a large allowed region within the global solar LMA contour. The inclusion of the KamLAND spectral data in the global fit gives a best-fit $\\Delta m^2 = 7.17 \\times 10^{-5}$ eV $^2$ and $\\tan^2\\theta=0.43 $ and constrains the allowed areas within LMA, leaving essentially two allowed zones. Maximal mixing though allowed by the \\kl data alone is disfavored by the global solar data and remains disallowed at about $3\\sigma$. The low $\\Delta m^2$ solution (LOW) is now ruled out at about 5$\\sigma$ with respect to the LMA solution. ",
        "introduction": "It is fair to say that the recently  announced first results  of  the Kamioka Liquid scintillator Anti-Neutrino Detector (KamLAND) experiment \\cite{kam_1} constitute a highly anticipated  milestone in our understanding and resolution of the three decade old solar neutrino problem. The origins of this puzzle lie in the  early deficit measurements of the solar neutrino flux in the pioneering Homestake chlorine experiment\\cite{cl}. This discrepancy between the expected rate, as predicted by increasingly refined solar model calculations \\cite{jnb_1} and the measured one has been  subsequently confirmed and buttressed over the years by results from the $^{71}{Ga}$ experiments SAGE, GALLEX and GNO \\cite{gno,sage}, the  Kamiokande and  the Super-Kamiokande experiments (SK) \\cite{superk}, and most recently from the Sudbury Neutrino Observatory (SNO) \\cite{sno1,sno2}. In particular, SK has provided valuable  zenith angle and energy spectrum information  in addition to total rate measurements of the high energy Boron flux, and SNO has provided crucial neutral current (NC) and charged current (CC) rate data along with spectrum  results. Over the years, these experimental results have been culled together with our understanding of neutrino mass, mixing and resonant matter oscillations  to obtain the allowed parameter space in terms of the mixing angle $\\tan^2\\theta$ and mass-squared difference $\\Delta m^2$  of the neutrino states. The analysis of global solar data carried out by various groups favours  the LMA solution based on  MSW resonant matter oscillations \\cite{msw} as the most probable resolution of the solar neutrino problem \\cite{a1,Bandyopadhyay:2002xj,Choubey:2002nc,Bandyopadhyay:2002qg, Bahcall:2002hv,a4,a5,a6,Strumia:2002rv,Fogli:2002pt,a9}.  KamLAND \\cite{kam_2} is a 1 kton liquid scintillator neutrino detector, designed specifically to test the LMA solution. It is located at the earlier Kamiokande site in the Kamioka mine in Japan. Its  main objective is to look for oscillation of $\\bar{\\nu_e}$ coming from Japanese nuclear power reactors situated at distances ranging from $\\sim$ $80$ km  to $800$ km. The bulk ($\\sim 79\\%$) of the measured flux is however from reactors which are at distances between 138 km to 214 km. The $\\bar\\nu_e$s are detected via the inverse beta decay reaction $\\bar{\\nu}_e+p\\rightarrow e^{+}+n$. Both the scintillation emitted by the positron as it moves through the detection medium, and its subsequent annihilation with an electron are recorded. The delayed coincidence of the positron with the  2.2 MeV $\\gamma$-ray from the capture of the neutron constitutes  a largely background free  signal. The total visible energy ($E_{vis}$) corresponds to $E_{e^+}+m_{e}$, where $E_{e^+}$ is the total energy of the positron and $m_e$ the electron mass. The positron energy is related to the incoming antineutrino energy as $E_{e^+}=E_{\\nu}- {\\bar E_{rec}} -(m_n -  m_p)$~MeV ($m_n - m_p = 1.293~$MeV is the neutron--proton mass difference). ${\\bar E_{rec}}$ is the average neutron recoil energy calculated here  using \\cite{beavo}. The energy resolution is $\\sigma(E)/E=7.5\\%/\\sqrt{E}$, $E$ is in MeV.  The first data from  KamLAND gives the ratio of the observed number of events to the expected number of events to be \\cite{kam_1} \\be R_{KL}=0.611 \\pm 0.085 (stat) \\pm 0.041 (syst) \\label{data} \\ee for an exposure of 162 ton-yr and a visible energy above 2.6 MeV \\footnote{Below this energy the background due to the geophysical neutrinos dominate.}. They have also presented the observed positron energy spectrum.  In a pre-\\kl  analysis \\cite{raj} we have shown that an energy integrated rate in the range 0.3 -0.8 will provide confirmation for the LMA solution. In particular the solar LMA best-fit predicted a \\kl rate of 0.65 which is close to the observed rate. This is  the first confirmation of the LMA solution to the solar neutrino problem using terrestrial neutrino sources. We also showed that for a rate below 0.9 the LOW solution to the solar neutrino problem is disallowed at more than 3$\\sigma$. Hence in this paper we focus on the LMA solution and perform a global analysis which combines\\\\ (i)KamLAND rate and  global solar data \\\\ (ii)KamLAND spectrum and global solar data\\\\ We  find the allowed area  from each of the above analyses and discuss  their contributions in sharpening our knowledge of  neutrino mass and mixing parameters. The current KamLAND and global solar data split the allowed LMA region in two parts -- a low $\\Delta m^2$ region (low-LMA) and a high $\\Delta m^2$ region (high-LMA), which has less (by $\\approx 2\\sigma$) statistical significance. A more precise determination of $\\Delta m^2$ and $\\tan^2\\theta$ should be possible with increased statistics and reduced systematics of the spectral data from KamLAND \\cite{bar,mur,deg,str,gg,aliani,Fogli:2002pb,raj}. We demonstrate the potential of 1 kton-yr spectral data in discriminating between the two allowed regions and further constraining the parameter values by simulating the spectrum at different values of $\\Delta m^2$ and $\\tan^2\\theta$ selected from the allowed area of the global solar+KamLAND  analysis. We find that if the true spectrum corresponds to that simulated at points in the low-LMA region then with 1 kton-yr exposure the high-LMA part can be further disfavoured. For spectrum simulated at high-LMA values however the ambiguity between the two zones persists. ",
        "conclusions": "In conclusion, we have investigated the impact of the first results from KamLAND on neutrino mass and mixing parameters in conjunction with the global solar neutrino data. KamLAND is completely consistent with the LMA solution, to the extent that the observed KamLAND rate is close to that predicted by the best-fit point of the LMA solution to the solar neutrino problem. As a result the combined analysis of the solar and KamLAND rates data allows a large area within the solar LMA region and the global solar best-fit does not change much with inclusion of \\kl rates. The constraining capabilities of the spectrum data is  much stronger and with only 145 days observed  spectrum \\kl can exclude certain parts of the LMA parameter space. After including the spectral data the allowed  LMA zone  consists mainly of two disconnected regions, one around the best-fit and another at a higher $\\Delta m^2$. The two zones merge at $3\\sigma$. Maximal mixing though allowed by the \\kl alone, is found to be still disfavored by the combined solar and \\kl data at more than $3\\sigma$. The LOW solution which was allowed at $3\\sigma$ from the global solar data and which predicts null oscillations in \\kl is now disfavored at almost $5\\sigma$ w.r.t the LMA solution.  With LMA now confirmed, the next focus of KamLAND would be a more accurate determination of the mass parameter by distinguishing between the two allowed sectors in the LMA region. We have explored this  through a projected analysis with 1 ktyr simulated data. With 1 kton-yr projected spectrum data the allowed $\\Delta m^2$ ranges around both low-LMA and high-LMA zones decrease in size and they get separated at 3$\\sigma$ by the spectrum data itself. The inclusion of the solar data disfavours(favours) the high(low)-LMA zone if the spectrum is simulated in the low(high)-LMA area. Thus the allowed areas become more precise for low-LMA spectrum while ambiguity between the two zones remains for high-LMA  spectrum. A higher statistics from \\kl is expected to resolve this ambiguity. A more precise determination of the mixing angle will be possible from a more accurate measurement of the CC/NC ratio at SNO.  \\vskip 10pt {\\bf  Acknowledgment } We acknowledge S. Pakvasa for useful discussions. RG would like to thank John Beacom for a helpful discussion. SC acknowledges discussions with S.T. Petcov."
    },
    "0212/astro-ph0212344.txt": {
        "abstract": "Phenomenological models for evolution of dusty galaxies and E/S0 galaxies, respectively, are developed to address two major questions concenring galaxy populations in deep infrared (IR) surveys: (1) Do normal late-type galaxies or starburst galaxies (including galaxies with obscured AGNs) dominate among sources in deep IR surveys? (2) How much do E/S0 galaxies contribute to the counts in deep mid-infrared (MIR: 3 -- 20$\\mu m$) surveys? Among three new models for evolution of dusty galaxies, it is assumed in Model S1 that starburst galaxies are the dominant population, and in Model S2 that normal galaxies dominate. Model S3 is an intermediate model. Comparing the model predictions with a wide range of observational data collected from the literature, we find that none of these models can be ruled out, given the uncertainties of the data. We show that the most direct method to distinguish these models is to compare the predicted color distributions of IR galaxies with observations, which will soon be available from the SWIRE survey. The models for E/S0 galaxies follow a simple passive evolution approach. Among the three E/S0 models (E1, E2 and E3) investigated in this paper, Model E2 which is specified by a peak formation redshift $z_{peak}=2$, and an e-folding formation time scale $\\omega=2$ Gyr, fits the data best. This suggests a synchronization between the evolution of E/S0 galaxies and of starburst galaxies, in the sense that the peak of the formation function of E/S0s ($z_{peak}=2$) is close to the peak of the evolution functions of starburst galaxies ($z_{peak} = 1.4$). We find that E/S0s contribute about 10 -- 30\\% of the counts in the MIR bands of $<10\\mu m$, and up to 30 -- 50\\% of the optical/NIR counts in the bright end. Their contributions to counts in the UV (2000{\\AA}) and in the longer wavelength IR ($\\geq 12\\mu m$) bands are negligible. Taking into account this contribution, new predictions for counts and confusion limits in the SIRTF bands are presented. ",
        "introduction": "Many new windows have recently been opened in various wavebands for observations of high redshift galaxies. Particularly, deep ISO surveys (e.g. Rowan-Robinson et al. 1997; Franceschini et al. 1997; Kawara et al. 1998; Elbaz et al. 1999; Aussel et al. 1999; Flores et al. 1999b; Puget et al. 1999; Dole et al. 2001; Clements et al. 1999; Serjeant et al. 2000) in the infrared and SCUBA surveys (Hughes et al. 1998; Barger et al. 1998; Blain et al. 1999) in the submm wavebands have shed lights on the dark side of galaxy formation and evolution. This stimulated a surging wave of empirical models (Xu et al. 1998; Blain et al. 1999; Roche \\& Eales 1999; Dole et al. 2001; Xu 2000; Rowan-Robinson 2001; Xu et al. 2001; Franceschini et al. 2001; Pearson 2001; Malkan \\& Stecker 2001; Chary \\& Elbaz 2001) to interpret the new ISO and SCUBA sources. The main aim of these empirical models is to better constrain some important characteristics of the galaxy evolution, including the peak epoch of the star formation rate in the universe and the roles played by different star formation modes (e.g. the quiescent mode against the interaction stimulated mode), for which the earlier optical deep surveys such as the Hubble Deep Field Survey (Williams et al. 1996) may have provided biased information (Madau et al. 1996) by neglecting the dust extinction (Rowan-Robinson  et al. 1997; Lonsdale 1999). This will certainly have impact on the theoretical simulations of galaxy formation and evolution, which (e.g. Somerville et al. 2001) have just started facing the fact that most of the early star formation may be hidden behind a thick veil of dust, making the incorporation of effects of dust extinction and emission in the framework essential (Granato et al. 2000).  Some of the most recent empirical models (Rowan-Robinson 2001; Xu et al. 2001; Franceschini et al. 2001; Pearson 2001; Chary \\& Elbaz 2001) have the feature that they can be constrained by counts and the cosmic background radiation in various IR and submm bands simultaneously. This is achieved by using SED templates to link sources in different bands. The SED library of Xu et al. (2001) contains realistic SEDs of a complete sample of 837 IRAS 25$\\mu m$ selected galaxies, enabling prediction of counts as well as color distributions, which provide additional constraints to the model. Templates in Rowan-Robinson (2001), Xu et al. (2001) and Chary \\& Elbaz (2001) extend from the IR-submm to optical and UV wavebands, hence these models can predict also the contributions from IR sources to optical and UV counts, relating the evolution of IR sources to that of galaxies seen in earlier optical and UV surveys. All of these works found strong evolution among IR sources in the redshift range of $0<z<1.5$. Assuming that the narrow sub-mJy bump on the Euclidian normalized differential counts of ISOCAM 15$\\mu m$ surveys (Elbaz et al. 1999) is due to the K-corrections caused by the strong unidentified broad band emission features (UIB) at 6 -- 8$\\mu m$, which are shifted into the ISOCAM 15$\\mu m$ band filter (LW3) when $z\\sim 1$, Xu (2000) concluded that typical IR galaxies at $z\\sim 1 $ have $L_{15\\mu m} \\sim 10^{11}\\; L_\\sun$, namely about 20 times more luminous than their local counterparts, while their comoving density is about the same as their local counterparts. As argued by Xu et al. (2001), the location of the 15$\\mu m$ bump ($f_{15\\mu m} \\sim 0.4 mJy$) provides a strong constraint on the luminosities of IR sources at $z\\sim 1$. This has been confirmed by the detailed study of ISOCAM sources in the HDF-north by Elbaz et al. (2002), who obtained redshifts for nearly all of these sources (40 out of total 41). The strong evolution has also been supported by surveys in other ISOCAM bands. For example, Clements et al. (2001) found that the luminosity functions of their ISO 12$\\mu m$ sources are consistent with a pure luminosity evolution of rate $\\sim (1+z)^{4.5}$, as derived by Xu (2000) from the ISOCAM 15$\\mu m$ counts.  On the other hand, different authors identified different populations of galaxies as the major carriers of the evolution of IR sources. It was found in early IRAS studies (Franceschini et al. 1988; Rowan-Robinson \\& Crawford 1989) that IR sources can be divided into 3 different populations: normal late-type galaxies ('cirrus galaxies'), starburst galaxies, and galaxies with AGNs. Assuming that all IR sources undergo pure luminosity evolution with the same evolution rate, Rowan-Robinson (2001) found that 'cirrus galaxies' dominate the ISOCAM $15\\mu m$ counts, the SCUBA 850$\\mu m$ counts, and the cosmic IR background radiation. Xu et al. (2001), motivated by results of optical identifications of ISO galaxies by Flores et al. (1999b) and Aussel et al. (1999) which show a larger percentage of these sources are in interacting systems, assumed that most of the evolution of the IR sources is due to starburst galaxies (defined by warm FIR colors: $f_{60\\mu m}/f_{100\\mu m} \\geq 0.5$). Obscured AGNs with relatively low $f_{25\\mu m}/f_{60\\mu m}$ ratios (due to the very high extinction affecting even the MIR fluxes) may be misclassified as starbursts in Xu et al. (2001), who classified galaxies with AGNs using the criterion $f_{25\\mu m}/f_{60\\mu m} \\geq 0.2$. Franceschini et al. (2001) also assumed that the strong evolution is confined to the starburst population. Their starburst galaxies include all 'active' galaxies not classified as Sey~1, i.e. including the Sey~2s and the LINERs, in Rush et al. (1993). Pearson (2001) found that a separate population of ultra-luminous galaxies (ULIRGs) with L$\\sim 10^{12} L_\\sun$, confined in a very narrow range of redshift centered at $z\\sim 1$, is mostly responsible for the strong evolution seen in the ISOCAM 15$\\mu m$ counts. Since all these models can fit the IR-submm counts and the CIB within uncertainty limits, there is apparently a degeneracy concerning the population of galaxies that carry most of the evolution of IR sources. The degeneracy still remains even when the new, detailed information of the 15 micron universe (source counts, redshift distributions, luminosity functions at different redshifts, etc.; Elbaz et al. 2002; Franceschini et al. 2001) is considered (see Section 2).  One of the central issues in the current agenda of hierarchical galaxy formation simulation studies concerns the roles played by two different star formation modes (eg. Somerville et al. 2001; Kauffmann et al. 2001), namely the quiescent star formation mode and the interaction-induced star formation mode. There are apparent links between the normal quiescent galaxies and the quiescent star formation (as in the Milky Way and M31) and between starburst galaxies and the interaction-induced star formation (as in the Antennae Galaxies and in M82). Therefore, the answer to the question whether the normal galaxies or starburst galaxies dominate the faint IR counts will have important impact on galaxy evolution theories.  Another deficiency of the current models for IR sources is the neglection of early type galaxies. Although these galaxies have little dust emission, therefore do not contribute significantly to IR counts at wavelengths longer than $\\sim 10\\mu m$, they are found as an important population in ISO 6.7$\\mu m$ counts (Flores et al. 1999b; Aussel et al. 1999; Serjeant et al. 2000). Future mid-IR surveys such as those planned for SIRTF IRAC cameras (4 bands centered at 3.6$\\mu m$, 4.5$\\mu m$, 5.8$\\mu m$ and 8$\\mu m$, respectively) will certainly detect many of these galaxies.  In this paper, we shall address the above two issues, namely (1) developing new models to investigate how to break the degeneracy concerning the major evolution population in IR counts (using SIRTF data in particular), and (2) modeling the evolution of E/S0 galaxies and predicting their contributions to the MIR counts. The plan of the paper is as follows: in Section 2 a set of new models for evolution of dusty galaxies are presented, while in Section 3 models for evolution of E/S0 galaxies are developed and compared to the observations; predictions for the UV, optical and NIR counts by a set of composite models, each consisting of a dusty galaxy evolution model and an E/S0 evolution model, are presented and compared with data in Section 4. The predictions for counts and confusion limits in SIRTF bands by the same models are presented in Section 5. In Section 6 we show how to use color distributions of ISO sources and future SIRTF sources to answer the question whether starburst galaxies or normal late-type galaxies are dominant in the IR sources. Section 7 is devoted to discussion. A summary is given in Section 8.\\footnote{More detailed results on predicitons of models presented in this paper are available on request to C.K.X..} The $\\Lambda$-Cosmology ($\\Omega_\\Lambda =0.7$, $\\Omega_m = 0.3$, H$_0$=75 km sec$^{-1}$ Mpc$^{-1}$) is assumed throughout the paper. ",
        "conclusions": "\\subsection{Evolution of SEDs of Starburst Galaxies} % As pointed out in Xu et al. (2001), the most important assumption in our models for the evolution of dusty galaxies is that high redshift (i.e. $z \\gsim 1$) star-forming galaxies have the same SEDs as their local counterparts when the luminosity is the same. Since our SEDs cover from the UV all the way through to the radio waveband, the validity of this assumption demands similarity between high redshift and local galaxies in many physical conditions, a requirement that seems too strong to be fulfilled in the strict sense. Particularly, the long term star formation history plays an important role in the optical and NIR emission of galaxies, and it is obvious that high redshift galaxies have very different long term star formation history (i.e. much younger) than that of local galaxies. Does this mean that our predictions for the optical/NIR flux densities of high redshift dusty galaxies are flawed and therefore unreliable? Our argument to dispute this suspicion is based on the fact that high redshift galaxies (particularly IR selected galaxies), already detected or to be detected in the future surveys, are almost exclusively high luminousity galaxies (less luminous galaxies with high redshift are too faint to be detected). As their local counterparts, these high redshift ULIRGs must host very powerful starbursts or AGNs as energy sources which contribute much of the emission even in the NIR bands ($\\sim 40\\%$ for the local ULIRGs, Surace et al. 2000; Scoville et al. 2000). Therefore, as far as the detectable IR sources are concerned, the difference in the underlying old population between the high redshift dusty starforming galaxies and their local counterparts does not affect very seriously our model predictions.  Recent multi-wavebands observations of SCUBA galaxies (Smail et al. 2002), Lyman-Break Galaxies (Adelberger \\& Steidel 2000), and ISOCAM galaxies (Flores et al. 1999a; Aussel et al. 1999; Cohen et al. 2000; Elbaz et al. 2002) are consistent with our assumption that high redshift star-forming galaxies have similar SEDs as their local counterparts, particularly when binned according to the luminosity (as stressed by Adelberger \\& Steidel 2000). On the other hand, some new observations hint at possible systematic differences between high redshift and local ULIRGs. The HST image of SCUBA galaxy SMM J14011+0252 (Ivison et al. 2001) shows that the star formation activity in that source is widely spread (up to a few kpc), a situation remarkably different from typical ULIRGs found in the local universe (Sanders \\& Mirabel 1998). Compared to centrally concentrated starbursts found in most of the local ULIRGs, galaxies with such widely distributed starbursts are expected to have less steep MIR slopes (i.e. smaller $f_{25\\mu m}/f_{12\\mu m}$ ratio) and cooler FIR color (i.e. larger $f_{100\\mu m}/f_{60\\mu m}$ ratios). This is due to the less intense radiation field (less warm dust emission in the 25$\\mu m$ and 60$\\mu m$ bands) and less dust opacity (less extinction for MIR fluxes). Indeed, Chapman et al. (2002) found that two sources detected both by the FIRBACK 170$\\mu m$ band survey (Puget et al. 1999; Dole et al. 2001) and by SCUBA (Scott et al. 2000), one at z=0.91 and the other at z=0.46, have significantly cooler dust temperatures ($T_{dust}\\sim 30$K) compared to $T_{dust}\\sim 50$K found for typical ULIRGs such as Arp220. They argue that this may indicate the starbursts in these systems are also extended. Given the small amount of information and the possible bias for cooler galaxies due to the sub-mm selection (Chapman et al. 2002), it is still too early to tell whether high-redshift ULIRGa are systematically more extended, and therefore have SEDs closer to less luminous local interacting galaxies (such as the Antennae galaxies) which are in earlier stages along the merger sequence. Future surveys like SWIRE will address these questions.   \\subsection{E/S0 evolution and starburst evolution} Comparisons between predictions E/S0 evolution models and observational data (Fig.11--12) indicate that the intermediate model (Model E2), which assumes a peak formation redshift of $z_{peak}=2$, is the favorite among the three models. This is in agreement with some previous works (Franceschini et al. 1998; Rodighiero et al. 2001; Im et al. 2002). This model can reproduce well the trend in the color-z plots and also fit the redshift distribution very well. However, particular in the B-K v.s. z plot, the data show significantly larger dispersion than the model predictions. There are two possible causes for this: \\begin{description} \\item{(1)} In our simple models, we have assumed that all of the E/S0s are formed through the same starburst procedure, which means their stellar populations have the same metallicity. It is known (e.g. Worthy 1994) that local E/S0s have different metallicity which is a major cause of the different M/L ratios and colors among these galaxies. Therefore, by neglecting these effects, our models leave some of the scatters in the color distributions unaccounted for. \\item{(2)} Some of the blue E/S0s in the data (Franceschini et al. 1998) may not be true E/S0s. According to Im et al. (2002), many of these 'blue interlopers' have strong, narrow emission lines, suggesting that they are low-mass starbursts rather than  massive star-forming E/S0s. \\end{description}  In fact, there is still a lack of consensus in the definition of E/S0s in deep surveys. Using a strict algorithm selecting the most symmetric and smooth galaxies, Im et al. (2002) found far less blue sources among their E/S0 sample than, e.g., Schade et al. (1999) and Menanteau (1999) whose samples were selected with less strict algorithms. A related uncertainty in our models is the choice on the exclusion of galaxies younger than 1 Gyr (Section 3.2). This choice is not entirely arbitrary: pushing the cut-off toward younger ages means more galaxies with high L/M ratios, which in turn results in too high model predictions for optical and NIR counts. This shows that the question of where to put the boundary between E/S0s and post-starbursts deserves more investigation.  If indeed E/S0s are formed through mergers (Toomre 1978; Kormendy \\& Sanders 1992), then their evolution is linked to the evolution of starburst galaxies which are closely related to mergers (Sanders \\& Mirabel 1998). Our results indeed indicate some synchronization between the two populations, in the sense that the peak of the formation function of E/S0s in the best fit model ($z_{peak}=2$) is close to the peak of the evolution functions of starburst galaxies ($z_{peak} = 1.4$). Future works exploring this possible link will provide important constraints to the evolution of both populations.  \\subsection{IR-quiet star-forming galaxies} % Our results (Fig. 13 -- 15) show that only a small fraction (10 -- 20 \\%) of UV selected galaxies (as in the sample of FOCA survey, Milliard et al. 1992) are IR bright, as indicated by the small contribution of simulated dusty galaxies to the UV counts (the E/S0 galaxies contribute even less). This suggests a separate population of IR quiet, UV bright galaxies which dominate the UV selected samples. Such a population is also needed for the B-band counts, because our simulations under-predict 30 -- 50\\% of observed counts in that band, while fully account for the R and K band counts (Fig. 13 -- 15). The best candidates for such galaxies are the low-metallicity (therefore low dust content) blue dwarf galaxies such as I Zw 18 (Searle \\& Sargent 1972). There is an apparent link between this IR quiet star-forming galaxy population and the 'faint blue galaxies' found in deep optical surveys (see Koo \\& Kron 1992 and Ellis 1997 for reviews). What is the relation between this population and the dusty (IR bright) star-forming galaxies? How does this population evolve (backwardly) with the redshift, and how does this evolution correlate with the evolution of IR bright galaxies? Are Lyman Break Galaxies, being selected by the UV flux in the rest frame, more closely related to the IR quiet star forming galaxies, or to the dusty star forming galaxies (as argued by Adelberger \\& Steidel 2000)? Answers to these questions will help to unify the pictures of galaxy formation/evolution seen in different wavebands. We plan to address these questions in our future work, particularly in connections with the forthcoming GALEX and SWIRE missions."
    },
    "0212/astro-ph0212229_arXiv.txt": {
        "abstract": "We report the most complete analysis to date of observations of the Cosmic Microwave Background (CMB) obtained during the 1998 flight of \\boomn.  We use two quite different methods to determine the angular power spectrum of the CMB in 20 bands centered at $\\ell = 50$ to 1000, applying them to $\\sim 50$\\% more data than has previously been analyzed.  The power spectra produced by the two methods are in good agreement with each other, and constitute the most sensitive measurements to date over the range $300 < \\ell < 1000$.  The increased precision of the power spectrum yields more precise determinations of several cosmological parameters than previous analyses of \\boom data.  The results continue to support an inflationary paradigm for the origin of the universe, being well fit by a $\\sim 13.5$~Gyr old, flat universe composed of approximately 5\\% baryonic matter, 30\\% cold dark matter, and 65\\% dark energy, with a spectral index of initial density perturbations $n_s \\sim 1$. \\\\ \\\\ ",
        "introduction": "Measurements of anisotropies in the cosmic microwave background (CMB) radiation now tightly constrain the nature and composition of our universe.  High signal-to-noise detections of primordial anisotropies have been made at angular scales ranging from the quadrupole~\\citep{bennett96} to as small as several arcminutes~\\citep{mason02,pearson02,dawson02}. The  power spectrum of temperature fluctuations shows a peak at spherical harmonic multipole $\\ell \\sim 200$ which has been detected with very high signal-to-noise by several teams~\\citep{debernardis00,hanany00,halverson01,scott02}, and strong indications of peaks at higher $\\ell$ have also been found~\\citep{halverson01,netterfield01,debernardis01}.  Within the context of models with adiabatic initial perturbations, as are generally predicted by inflation, these measurements have been used in combination with various other cosmological constraints to estimate the values of many important cosmological parameters.  Combining their CMB data with weak cosmological constraints such as a very loose prior on the Hubble constant, various teams have made robust determinations of several parameters, including the total energy density of the universe $\\Omega_{total}$, the density of baryons $\\Omega_b$, and the value of the density perturbation power spectral index, $n_s$ ~\\citep{lange01,balbi00,pryke01,netterfield01}. Many other parameters are tightly constrained when stronger constraints on cosmology are assumed.  We report here new results from the 1998 Antarctic flight of the \\boom experiment.  Previous results from this flight using less data than included here were published in \\cite{debernardis00} (hereafter B00) and \\cite{netterfield01} (hereafter B02). Here we use the two very different analysis methods of B00 and B02, and apply them  over a larger fraction of the dataset to make an improved measurement of the CMB angular power spectrum. ",
        "conclusions": "In this paper we have presented an analysis of 50\\% more data from the 1998 Antarctic flight of \\boom than previously treated. Our analysis is the most thorough to date, using two very different power spectrum estimation pipelines to derive the angular power spectrum of the cosmic microwave background radiation.  The two methods show good agreement and, with the greater amount of data used, an increase in the precision of measured power spectrum.  In particular, features in the power spectrum beyond the first peak (at $\\ell \\sim 200$) are detected with greater confidence.  Given that such features are a natural consequence of standard cold dark matter dominated cosmological models with adiabatic initial density perturbations, their presence gives us greater confidence in the validity of that model set.  Within the context of these models we have estimated the the values of cosmological parameters using the results from both of our analysis methods.  The resulting parameter values are insensitive to the small differences between the two results. At the increased precision with which we determine the cosmological parameters, we find that our results remain completely consistent with a flat $\\Lambda$-CDM cosmology."
    },
    "0212/astro-ph0212535_arXiv.txt": {
        "abstract": "We revisit the process of gravitational sedimentation of helium and heavy elements in the intra-cluster medium.  We find that helium applies an inward drag force on heavy elements, boosting their sedimentation speed to nearly half its own.  This speed is almost independent of the mass and the electric charge of heavy elements. In the absence of  small-scale magnetic fields, helium sedimentation can increase the He/H abundance ratio in the cores of hot clusters by three orders of magnitude.  It also steepens the baryonic density profile yielding a higher X-ray luminosity, which offers an explanation of the observed luminosity-temperature relation.  If the primordial He/H ratio is assumed, then the gas density inferred from the observed X-ray emissivity might be underestimated by $ 30\\%$ in the cores of clusters and overestimated by $7\\%$ in the outer regions.  The dark matter density on the other hand might be overestimated by a factor of 8/3 in the cores and underestimated by $ 18 \\%$ in outer regions. ",
        "introduction": "We write the equations governing the evolution of individual species in an ICM of any composition. We do not consider magnetic fields in this paper, but include electric fields which must exist in any ionized plasma in a gravitational field (Eddington 1926).  Let the ICM be made of any number of species each made of particles characterized by mass $m_{\\rm i}$ and electric charge $q_{\\rm i}=e Z_{\\rm i}$. We denote the local number density and velocity of a patch of matter of each species by $n_{\\rm i}$ and $\\vV_{\\rm i}$. Then, the mass density is $\\rho_{\\rm i}=n_{\\rm i} m_{\\rm i}$ and the pressure is $P_{\\rm i}=n_{\\rm i} kT$, where we have assumed that the ICM is in local thermal equilibrium so all species share the same temperature $T$. With this notation the equations are \\begin{eqnarray} \\label{h2} \\frac{\\pa \\vV_{\\rm i}}{\\pa t} +\\vV_{\\rm i}\\cdot \\vnabla \\vV_{\\rm i} =-\\frac{\\vnabla P_{\\rm i}}{\\rho_{\\rm i}}+\\vg+\\frac{q_{\\rm i}\\vE}{m_{\\rm i}}+\\sum_{j}(\\vV_{\\rm i}-\\vV_{\\rm j})/\\tau_{\\rm ij} \\, , \\end{eqnarray} where $\\vE$ is the electric field, and $\\vg$ is the gravitational field (force per unit mass). The constants $\\tau_{\\rm ij}$ are the time-scales for the drag forces acting on the species $i$ as a result of Coulomb interactions with species $j$.  Momentum conservation implies that $\\tau_{\\rm ij}=(\\rho_{\\rm j} /\\rho_{\\rm i})\\tau_{\\rm ji}$.  If $m_{\\rm i}\\gg m_{\\rm j}$ then $\\tau_{\\rm ij}$ can be approximated by (Spitzer 1968) \\begin{eqnarray} \\label{Sp} \\tau_{\\rm ij}&=&\\frac{3(2\\pi)^{1/2}(kT)^{3/2}}{8\\pi Z_{\\rm i}^2Z_{\\rm j}^2e^4n_{\\rm j} \\ln \\Lambda}\\frac{m_{\\rm i}}{m_{\\rm j}^{1/2}}\\\\& =& \\nonumber 9.4\\times 10^6\\left(\\frac{T}{10^8 {\\rm K}}\\right)^{3/2} \\left(\\frac{{\\rm \\ln}\\Lambda}{40}\\right)^{-1} \\left(\\frac{n_{\\rm j}}{10^3 {\\rm m}^{-3}}\\right)^{-1}\\\\&\\nonumber\\times&\\left(\\frac{m_{\\rm i}}{m_{\\rm p}}\\right) \\left(\\frac{m_{\\rm j} }{m_{\\rm p}}\\right)^{-1/2}\\left(Z_{\\rm i} Z_{\\rm j}\\right)^{-2}\\, {\\rm Yr}\\, , \\end{eqnarray} where the ${\\rm ln} \\Lambda$ is the Coulomb logarithm. The equations (\\ref{h2}) must be  supplied by additional relations that specify  the electric field. To create an electric force comparable with gravitational and pressure forces a tiny fractional  charge excess (${Gm_{\\rm p}^2}/{e^2}\\sim 10^{-36}$) is sufficient. The additional relations can then be obtained by assuming charge neutrality and zero  electric currents, i.e., \\begin{equation} \\sum_{i} n_{\\rm i} q_{\\rm i}=0 \\, ,\\quad {\\rm and} \\quad \\sum_{\\rm i} n_{\\rm i} q_{\\rm i}\\vV_{\\rm i}=0\\, . \\label{h5} \\end{equation}  \\subsection{Sedimentation speeds and time-scales}   Multiplying the equations (\\ref{h2}) by $n_{\\rm i} m_{\\rm i}/\\sum n_{\\rm i} m_{\\rm i} $ and summing over all species yields \\begin{eqnarray} \\frac{\\vnabla P}{\\rho}= \\tvg \\, , \\label{h1} \\end{eqnarray} where $\\rho=\\sum n_{\\rm i} m_{\\rm i}$ is the total local density of all species, $\\vV=\\sum \\rho_i \\vV_{\\rm i} /\\rho$ and $P=\\sum P_{\\rm i}$ are, respectively, the mass-weighted mean velocity and total pressure, and $\\tvg=\\vg-\\pa \\vV/\\pa t - \\vV \\cdot \\vnabla \\vV$ is the gravitational field in the frame of reference moving with the velocity $\\vV$.  The terms involving the electric and drag forces have disappeared by virtue of charge neutrality and momentum conservation, respectively. Since $|\\vV_{\\rm i}-\\vV|$ is much smaller than the thermal velocities, the velocity dispersion term $\\sum \\rho_i (\\vV_{\\rm i}-\\vV)\\vnabla (\\vV_{\\rm i}-\\vV)/\\rho$ is negligible compared to the pressure and gravity terms we have not included it in the equation (\\ref{h1}).  At the initial stage of the diffusion process the dominant light species have the same distribution, which gives ${\\grad (n_{\\rm i} kT)}/{(n_{\\rm i} m_{\\rm i})}=(\\mu m_{\\rm p}/m_{\\rm i})\\tvg$, where $\\mu=(1/m_{\\rm p}) \\sum n_{\\rm i} m_{\\rm i}/\\sum n_{\\rm i}$ is the mean molecular weight. Thus equations (\\ref{h2}) yield \\begin{eqnarray} \\label{h3} (\\mu m_{\\rm p}/m_{\\rm i}-1)\\tvg+\\frac{q_{\\rm i} \\vE}{m_{\\rm i}}+\\sum_{j}(V_{\\rm i}-V_{\\rm j})/\\tau_{\\rm ij }=0\\; . \\end{eqnarray} Using equations (\\ref{h5}) and (\\ref{h3}) we find that the electric field is $\\vE=\\mu \\tvg/e$. From (\\ref{Sp}) $\\tau_{\\rm ij}\\propto n_{\\rm j}^{-1}m^{-1/2}_{\\rm j}$, so because of  the  low abundance of heavy ions and the low mass of electrons, only the drag terms from  protons and helium ions is important. So we are left with a single equation for each element \\begin{eqnarray} \\label{hs} \\left(\\frac{1+Z_{\\rm i}}{A_{\\rm i}}\\; \\mu-1\\right)\\tvg+ \\frac{\\vV_{\\rm p}-\\vV_{\\rm i}}{\\tau_{\\rm ip}}+\\frac{\\vV_{\\rm \\alpha}-\\vV_{\\rm i}}{\\tau_{\\rm i\\alpha}}=0 \\; . \\end{eqnarray}  We can now estimate the velocity of each species relative to $\\vV_{\\rm p}$, the velocity of the proton fluid at each point. The advantage of using velocities relative to protons is the independence of the results from physical processes other than diffusion (radiative cooling, heating by supernovae etc). Even though these processes can have important dynamical effects on the cluster, they do not affect the expressions we develop for the relative velocities and element abundance.  From (\\ref{hs}) the helium velocity is \\begin{eqnarray} \\vV_\\alpha-\\vV_{\\rm p} =(3\\mu/4-1)\\tvg\\; \\tau_{\\rm \\alpha p}\\approx -0.56\\tvg\\; \\tau_{\\rm \\alpha p}. \\end{eqnarray}  The relative sedimentation speed of heavy elements, $\\vV_{\\rm i}-\\vV_{\\rm p}$, can easily be related to that of helium $\\vV_\\alpha-\\vV_{\\rm p}$.  Heavy elements experience an upward proton drag and inward helium drag forces. According to (\\ref{Sp}), $\\tau_{\\rm ij}\\propto Z_{\\rm i}^{-2}$ and for large $Z_{\\rm i}$ these drag terms dominate all others in (\\ref{hs}).  Thus we are left with \\begin{eqnarray} \\label{vel} (\\vV_p-\\vV_{\\rm i})/\\tau_{\\rm ip}\\approx (\\vV_{\\rm i}-\\vV_{\\alpha})/\\tau_{\\rm i\\alpha}. \\end{eqnarray} Substituting $\\tau_{\\rm i\\alpha}$ and $\\tau_{\\rm i p}$ from (\\ref{Sp}) in the last equation yields \\begin{eqnarray} \\frac{\\vV_{\\rm i}-\\vV_{\\rm p}}{\\vV_{\\alpha}-\\vV_{\\rm p}}\\approx \\left(1+\\frac{n_{\\rm p}}{8n_\\alpha }\\right)^{-1}= 0.4\\, , \\end{eqnarray} where we have taken $n_\\alpha/n_{\\rm p}=0.08$. Thus, contrary to the prediction of Fabian \\& Pringle (1977), who by neglecting helium drag obtained $(V_{{\\rm p}}-V_{\\rm i})\\propto A_{\\rm i} Z_{\\rm i}^{-2}$, the diffusion speeds are approximately the same for all heavy elements. Including all forces in the calculation increases slightly this result (by a factor of $\\sim 1+4/Z_{\\rm i}$) \\begin{eqnarray} \\frac{\\vV_{\\rm i}-\\vV_{\\rm p}}{\\vV_{\\alpha}-\\vV_{\\rm p}}=0.4 \\left[1+\\frac{3A_{\\rm i}}{Z_{\\rm i}^2}-1.8\\left(\\frac{1}{Z_{\\rm i}}+\\frac{1}{Z_{\\rm i}^2}\\right)\\right]. \\label{ratspeed} \\end{eqnarray} There is also  a correction of a similar magnitude  if the initial distribution of heavy elements is different from that of light  elements.  The relation (\\ref{ratspeed}) is independent of the physical state of the ICM.  To obtain the relative velocities we have to know the temperature, the density and the gravitational acceleration.  Taking in (\\ref{Sp}) ${\\rm ln} \\Lambda\\approx 40$ as the typical value for the ICM, we find the helium velocity relative to protons to be \\begin{equation} \\vV_{\\alpha}-\\vV_{\\rm p}=5\\times 10^4 \\tvg_{_{-9.5}}{n_{\\rm p}^{-1}}_{_{-3}} T_{_8} \\hspace{3mm} {\\rm m s^{-1}}, \\label{speeds} \\end{equation} where $\\tvg_{_{-9.5}}=\\tvg/(10^{-9.5} {\\rm ms^{-2}})$, ${n_{\\rm p}}_{_{-3}} =n_{\\rm p}/(10^3 {\\rm m^{-3}})$, and $T_{_8}=T/10^8$ K. This is lower by $\\sim 40\\%$ than the result of Qin \\& Wu (2000), the main difference is the inclusion of electric field in our calculation. As seen from equation (\\ref{speeds}) the diffusion speeds depend on the local density and temperature. However, a single diffusion time-scale is obtained if the gas and dark matter both follow an isothermal spherical density profile $\\rho(R)\\propto R^{-2}$ and the gas is in hydrostatic equilibrium ($\\tilde g=g$) with constant $T$. In this case $V_{\\rm i}-V_{\\rm j}\\propto g/n_{\\rm p}\\propto R$, which together with continuity equation, $\\dd n_{\\rm i}/\\dd t=-R^{-2} \\partial_R(R^2 V_{\\rm i})$, yields, \\begin{eqnarray} \\frac{\\rm d}{{\\rm d}t}\\ln\\left(\\frac{n_{\\rm i}}{n_{\\rm j}}\\right)= \\frac{3(V_{\\rm j}-V_{\\rm i})}{R} \\; . \\label{cont} \\end{eqnarray} This motivates us to define the time-scale for diffusion between two species as \\begin{equation} \\tau_D=\\frac{R}{3(V_{\\rm i} -V_{\\rm j} )}, \\end{equation} which for helium relative to  protons gives \\begin{eqnarray} \\tau_D= 3\\times 10^9 \\left(\\frac{f_{\\rm b}}{0.1}\\right) \\left(\\frac{T}{10^8 {\\rm K}}\\right)^{3/2} {\\rm Yr}\\; , \\label{heliumtau} \\end{eqnarray} where $f_{\\rm b}$ is the baryonic mass  fraction in the cluster. ",
        "conclusions": "We have seen that diffusion steepens the gas density profile near the center of hot clusters, increasing their X-ray luminosity (by factor of 5 in our example).  In colder clusters ($T\\sim 10^7$ K) the diffusion time-scale is larger than the Hubble time and their luminosity remains unchanged. This may explain the observed discrepancy between the observed $L\\propto T^{3}$ (Mushotzky 1984; Edge \\& Stuwart 1991; David et al. 1993) and $L\\propto T^{2}$ which is expected from self-similar arguments (e.g., Kaiser 1986).   If the inner regions of clusters are dominated by helium then the baryonic mass density as inferred from the X-ray emissivity can be underestimated by $\\sim 30\\%$ if constant He/H abundance ratio is assumed, while in the helium-deficient outer region it can be overestimated by $\\sim 7\\%$.  Estimates of of the dark matter density can be affected by even a larger factor. These estimates assume hydrostatic equilibrium (eq. \\ref{hstat}), so by taking $\\mu= 0.59$ (cosmic abundance) instead of 0.5 (pure hydrogen plasma) in the outer regions, we underestimate the total mass by $\\sim 18 \\%$ (Qin \\& Wu 2000).  In the helium dominated core the mass would be overestimated by a factor of $2.3$.  Our estimates of the sedimentation time-scales have neglected magnetic fields.  Magnetic fields with coherence length comparable to the size  of the cluster force the ions to move on longer orbits defined by the field lines.  This can increase the sedimentation time-scales by factor of a few.  Small-scale magnetic fields, however,    can increase the time-scales by a factor ranging from a few in some estimates (Narayan \\& Medvedev 2001; Malyshkin 2001) to 100-1000 in others (Chandran \\& Cowley 1998)."
    },
    "0212/astro-ph0212159_arXiv.txt": {
        "abstract": "We present high-resolution ($R$ = 20000) spectroscopy of H$_3^+$ absorption toward the luminous Galactic center sources GCS~3-2 and GC IRS~3.  With the efficient wavelength coverage afforded by Subaru IRCS, six absorption lines of H$_3^+$ have been detected in each source from 3.5 to 4.0~$\\mu$m, three of which are new. In particular the 3.543~$\\mu$m absorption line of the $R$(3, 3)$^l$ transition arising from the metastable ($J$, $K$) = (3, 3) state has been tentatively detected for the first time in the interstellar medium, where previous observations of H$_3^+$ had been limited to absorption lines from the lowest levels: ($J$, $K$) = (1, 0) of ortho-H$_3^+$ and (1, 1) of para-H$_3^+$.  The H$_3^+$ absorption toward the Galactic center takes place in dense and diffuse clouds along the line of sight as well as the molecular complex close to the Galactic nucleus. At least four kinematic components are found in the H$_3^+$ absorption lines. We suggest identifications of the velocity components with those of H {\\sc i}, CO, and H$_2$CO previously reported from radio and infrared observations.  H$_3^+$ components with velocities that match those of weak and sharp CO and H$_2$CO lines are attributed to diffuse clouds. Our observation has revealed a striking difference between the absorption profiles of H$_3^+$ and CO, demonstrating that the spectroscopy of H$_3^+$ provides information complementary to that obtained from CO spectroscopy.  The tentative detection of the $R$(3, 3)$^l$ line and the non-detection of spectral lines from other $J > 1$ levels provide observational evidence for the metastability of the (3, 3) level, which is theoretically expected. This suggests that other metastable $J$ = $K$ levels with higher $J$ may also be populated. ",
        "introduction": "The crucial role which the H$_3^+$ molecular ion plays in interstellar chemistry was first addressed by \\citet{wat73} and \\citet{her73}. H$_3^+$ is produced by cosmic-ray ionization of H$_2$ to H$_3^+$ followed by the efficient Langevin reaction, \\begin{equation} {\\rm H_2^+ + H_2 \\rightarrow H_3^+ + H.} \\end{equation} It works as a universal protonator (acid) in the efficient proton hop reaction \\begin{equation} {\\rm H_3^+ + X \\rightarrow H_2 + XH^+ } \\end{equation} for most molecules or atoms X (He, Ne, N, and O$_2$ are a few exceptions). Subsequent general reactions with a molecule {\\rm Y}, \\begin{equation} {\\rm XH^+ + Y \\rightarrow XY^+ + H}, \\end{equation} lead further to chemical networks which produce complex molecules in dense interstellar clouds.  \\newpage A search for interstellar H$_3^+$ was initiated by \\citet{oka81} more than two decades ago, and culminated in the discovery of H$_3^+$ in the interstellar medium toward AFGL~2136 and W33A, young stellar objects deeply embedded in molecular clouds \\citep{geb96}. Since then the study of interstellar H$_3^+$ has progressed rapidly by a series of successful observations in dense clouds \\citep{mcc99} and in diffuse clouds toward the highly reddened star Cygnus OB2 12, the Galactic center, and many other sightlines \\citep{mcc98b, geb99, mcc02}. These observations have not only demonstrated the ubiquity and abundance of this molecular ion, which was anticipated in general in molecular hydrogen dominated plasmas \\citep{mar61}, but also revealed an interesting enigma related to the H$_3^+$ chemistry of the diffuse interstellar medium. While the H$_3^+$ column densities measured in dense clouds agreed well with those predicted by the H$_3^+$ production and destruction mechanisms of \\citet{her73} \\citep{geb96, mcc99}, those measured in the diffuse interstellar medium have been found to be orders of magnitude higher than estimated using canonical values of plasma chemical constants (McCall et al. 1998a,b; Geballe et al. 1999, McCall et al. 2002). In order to have a better understanding of this problem, we undertook an absorption-line survey of this molecular ion toward the Galactic center.  This study also provides unique information on the medium toward the Galactic center because of the special characteristics of H$_3^+$ as an astrophysical probe.  The absorption-line survey was designed to cover the maximum number of H$_3^+$ absorption lines in clouds with various physical conditions. The Galactic center sources are ideal for this purpose, since they suffer heavy visual extinction of $A_{V}$ = 25--40 \\citep{cot00}, the highest in the Galaxy among those obscured mainly by diffuse clouds. The line of sight to the Galactic center cuts through dense and diffuse clouds in the intervening spiral arms as well as the molecular complex close to the nucleus.  Thus, H$_3^+$ in different physical conditions is efficiently sampled in one observation.  The intervening clouds have different radial velocities, and the nature of the clouds were studied in previous spectroscopic observations using molecular and atomic tracers from near-infrared to radio wavelengths. Two infrared sources, GCS~3-2 in the Quintuplet cluster \\citep{nag90,oku90} and GC~IRS~3 \\citep{bec75, bec78} near Sgr A$^\\ast$ have been selected, since both appear to have intrinsically featureless spectra due to dust emission and are sufficiently luminous to provide good continuum fluxes for $L$-band absorption spectroscopy.  Wide wavelength coverage is essential for a line survey. The observation of interstellar H$_3^+$ has so far been almost entirely limited to the three absorption lines, [$R$(1, 0), $R$(1, 1)$^l$, and $R$(1, 1)$^u$], all arising from the lowest $J$ = 1 rotational levels, i.e.\\ ($J$, $K$) = (1, 0) of ortho-H$_3^+$ and (1, 1) of para-H$_3^+$ (note that the $J$ = 0 level is forbidden by the Pauli principle; McCall 2001). This is because high spectral resolution and wide wavelength coverage are often incompatible. This limitation has now been overcome by a new generation of cross-dispersed infrared spectrographs, such as IRCS on the Subaru Telescope.  We can now simultaneously observe at high spectral resolution not only most of the transitions starting from the $J$ = 1 levels, including $Q$(1, 0) and $Q$(1, 1), but also other transitions starting from higher $J$ levels. This wide coverage has led us to a tentative detection of the $R$(3, 3)$^l$ line starting from the metastable ($J$, $K$) = (3, 3) level. ",
        "conclusions": "\\subsection{Location of H$_3^+$ Absorbing Clouds} We discuss below possible identifications of the absorbing clouds, referring to previous observations at radio and infrared wavelengths s(figure \\ref{fig:h2co}). We will base the following velocity discussion on GC~IRS~3 spectral lines unless otherwise stated, since the line profiles are sharper than those of GCS~3-2.  \\subsubsection{$-$140~km~s$^{-1}$ to $-$110~km Component}  The $-$140~km~s$^{-1}$ component is attributed to clouds in the ``expanding molecular ring'' \\citep{kai72,sco72}. The ``expanding molecular ring'' is a chain of molecular clouds orbiting around the nucleus at 200~pc from the Galactic center, and gradually receding from it. The $-$140~km~s$^{-1}$ component of GC~IRS~3 appears at $-$110~km~s$^{-1}$ in GCS~3-2, which is reasonable because the ``expanding molecular ring'' appears with a less negative velocity at positive Galactic longitude. The additional possible peak at $-$170~km~s$^{-1}$ seen towards GCS~3-2 could be the high velocity component discussed by \\citet{gue81}, \\citet{lis93}, and \\citet{yus93}.  \\subsubsection{$-$60~km~s$^{-1}$ Component}  A clear match of the peak velocities at $-$60~km~s$^{-1}$ is noted between the H$_3^+$ lines and the H$_2$CO line observed in absorption toward Sgr A (Snyder et al. 1969; G\\\"usten, Downes 1981), and the H {\\sc i} absorption line \\citep{lis85}. The cloud velocity is consistent with the radial velocity of the ``3~kpc arm''. The ``nuclear disk'' rotating at 300~pc from the Galactic center (Sanders, Wrixon 1973) also shows a similar velocity at the longitude of GC~IRS~3 (Whiteoak, Gardner 1979). However, we infer that the contribution of the ``3~kpc arm'' clouds is dominant, since the component is also seen in GCS~3-2 at the same velocity. This interpretation is also supported by the absence of a positive velocity component which is expected from the ``nuclear disk'' for the positive Galactic longitude of GCS 3-2.  \\subsubsection{0~km$^{-1}$ Component}  The 0~km~s$^{-1}$ component is also apparent in the H$_2$CO and H {\\sc i} absorption spectra of Sgr A in figure \\ref{fig:h2co}. The 0~km~s$^{-1}$ absorbers are usually attributed to the ``local clouds'' within a few kpc of the solar neighborhood, which is most evident in the H {\\sc i} absorption at 21~cm (e.g., Garwood, Dickey 1989). There should also be a contribution from the low-velocity dense clouds close to Sgr~A. Some of them are reported to be within about 50~pc of Sgr~A by \\citet{whi78}.  \\subsubsection{+50~km$^{-1}$ Wing} A candidate for the H$_3^+$ absorption wing at +50~km~s$^{-1}$ is the ``+50~km~s$^{-1}$ clouds,'' a complex of giant molecular clouds within about 10~pc of the Galactic nucleus (G\\\"usten et al. 1981; G\\\"usten, Henkel 1983). In the model of the large-scale structures in the central 10~pc of our Galaxy constructed with recent high spatial resolution radio observations (Coil, Ho 2000; Wright et al. 2001), the Sgr~A~East non-thermal radio source is impacting the ``+50~km~s$^{-1}$ clouds'' at the far side of the Galactic center. If the H$_3^+$ absorption wing at +50~km~s$^{-1}$ is associated with these giant molecular clouds, it places GC~IRS~3 beyond Sgr A, which then jeopardizes the membership of GC~IRS~3 in the central star cluster. Instead, to account for the absorption feature of CO in the near-infrared at 4.7~$\\mu$m, \\citet{geb89} proposed that the particular cloud occulting GC~IRS~3 could be a part of the ``circumnuclear disk'' delineated by \\citet{gue87} in their HCN map. The ``circumnuclear disk'' is a compact ($\\sim$ 3~pc) clumpy torus, a reservoir supposedly feeding mass to the dynamical center of our Galaxy, Sgr~A$^\\ast$. The +50~km~s$^{-1}$ component is barely seen in GCS~3-2, which suggests that the absorber of GC~IRS~3 is a compact local structure at Sgr~A, and thus argues in favor of the ``circumnuclear disk'' origin.  \\subsection{Special Characteristic of the H$_3^+$ Number Density}  Lines of sight toward the Galactic center sample both dense and diffuse clouds along the long pathlength. About 30\\% of the visual extinction toward the Galactic center is believed to arise in dense clouds, and the rest in diffuse clouds (e.g. Whittet et al. 1997). For instance, the diffuse and dense cloud extinctions to GC~IRS~3 estimated from the optical depths of hydrocarbon and water ice in the 3~$\\mu$m region \\citep{chi02} are $A_{V}$ = 36 and $A_{V}$ = 11, respectively.  The optical depth of H$_3^+$, however, does not scale with the visible extinction. In contrast to most other molecules, the number density of H$_3^+$ is constant and independent of the density of the cloud as long as the number density of H$_2$ relative to that of the destroyer X of H$_3^+$, $n$(H$_2$)/$n$(X), is constant (figure \\ref{fig:dens1}). A simple analysis using a steady state chemical kinetics yields a H$_3^+$ number density of $n$(H$_3^+$) = ($\\zeta$/$k_{\\rm CO}$)[$n$(H$_2$)/$n$(CO)] for a typical dense cloud where carbon atoms are mostly in the form of CO, and $n$(H$_3^+$) = ($\\zeta$/$k_{\\rm e}$)[$n$(H$_2$)/$n$(${\\rm e}$)] for a typical diffuse cloud where carbon atoms are in the form of C$^+$ ($\\zeta$ is the cosmic ray ionization rate, and $k_{\\rm CO}$ and $k_{\\rm e}$ are the rate constants of the ion-neutral reaction between H$_3^+$ and CO and of dissociative electron recombination of H$_3^+$, respectively; McCall et al. 1998a,b). If typical values of $\\zeta \\sim 3 \\times 10^{-17}$ s$^{-1}$, $k_{\\rm CO} \\sim$ 2 $\\times$ 10$^{-9}$ cm$^3$ s$^{-1}$, $k_{\\rm e} \\sim 5 \\times 10^{-7}$ cm$^{3}$ s$^{-1}$ are used, and [$n$(H$_2$)/$n$(CO)] and [$n$(H$_2$)/$n$(${\\rm e}$)] are both assumed to be 7 $\\times$ 10$^3$, we obtain $n$(H$_3^+$) $\\sim$ 1 $\\times$ 10$^{-4}$ cm$^{-3}$ in dense clouds and $\\sim$ 4 $\\times$ 10$^{-7}$ cm$^{-3}$ in diffuse clouds. These constancies of the number density make H$_3^+$ nearly a direct indicator of the dimension of the absorbing clouds. In other words, regardless of how high or low the cloud density is, the H$_3^+$ column density is simply proportional to the physical dimension of the cloud along the line of sight. This simple analysis well explains the observed H$_3^+$ column densities in dense clouds \\citep{mcc99}.  However, a problem appears when this model is applied to diffuse clouds because of uncertainties of the constants used in the calculation. Since the observed H$_3^+$ column densities in dense and diffuse clouds are comparable, the large factor of 250 difference in $n$(H$_3^+$) leads to the same large difference in the cloud dimension, which is difficult to accept (McCall et al. 1998a,b; Geballe et al. 1999; McCall et al. 2002). Currently, at least one of the three assumed values, [$\\zeta$, $k_{\\rm e}$, and $n$(${\\rm e}$)] is under suspicion.  For instance, measurements of $k_{\\rm e}$ vary by more than one order depending on the experimental techniques (e.g. Larsson 2000). It is likely that the estimate of $n$(H$_3^+$) in diffuse clouds will be increased by one order of magnitude or so when the uncertainty in these three parameters is resolved; $n$(H$_3^+$) will then give a reasonable cloud dimension.  \\subsection{H$_3^+$ in Diffuse Clouds}  The sharp and weak H$_2$CO lines toward Sgr~A at $-$140~km~s$^{-1}$, $-$60~km~s$^{-1}$ and 0~km~s$^{-1}$, and the weakness of the CO absorption at these velocities in GC~IRS~3 indicate that the clouds responsible for those absorptions are diffuse. \\citet{geb89} estimated the CO column density in the cloud at $-$60~km~s$^{-1}$ velocity to be 1 $\\times$ 10$^{17}$ cm$^{-2}$ with a cloud temperature of 17~K. If we use [H$_2$]/[CO] = 7 $\\times$ 10$^3$, we obtain $N$(H$_2$) = 7 $\\times$ 10$^{20}$ cm$^{-2}$, though this value might be better taken as a lower limit because of the possible saturation in the CO fundamental lines.  Such a low hydrogen column density is consistent with diffuse clouds.  The observed H$_3^+$ column density in each of these velocity components of GC IRS 3 is in the range $N$(H$_3^+$) = (2.4--5.6) $\\times$ 10$^{14}$ cm$^{-2}$. If the H$_3^+$ number density in diffuse clouds given above (4 $\\times$ 10$^{-7}$ cm$^{-3}$) is assumed, we obtain unreasonably high cloud dimensions of 200 to 500~pc. The enigma of H$_3^+$ chemistry in the diffuse interstellar medium needs to be solved before we discuss the cloud dimension any further.  The large pedestal component of GC IRS 3 is also interpreted to be due to diffuse clouds, because of the absence of strong CO or H$_2$CO absorption. A clarification of the nature of this cloud is an interesting future problem. There exists a possibility that the decomposition of the observed features into velocity components is not unique in view of the significant noise in the spectrum and the limited spectral resolution.  \\subsection{H$_3^+$ in Dense Clouds}  The strong and saturated CO absorption toward GC~IRS~3 and H$_2$CO for Sgr~A at the velocity of +50~km~s$^{-1}$ wing clearly represents dense clouds (figure \\ref{fig:h2co}). It is remarkable that the H$_3^+$ absorption lines starting from the $J$ = 1 levels are weak at this velocity (figure \\ref{fig:spe1}). This must be due to a small pathlength of a very dense cloud in front of the infrared source. From the equivalent width of the $R$(1, 1)$^l$ line integrated over +20 to +60~km~$^{-1}$, $N$(H$_3^+$)$_{J=1}$ = 2.4 $\\times 10^{14}$ cm$^{-1}$ is obtained; this corresponds to a dense cloud pathlength of 1.1~pc if $n$(H$_3^+$) $\\sim 1 \\times 10^{-4}$ cm$^{-3}$ estimated in subsection 5.2 is assumed. The concentration of the H$_3^+$ in the (3, 3) level at the velocity would increase the total column density and the cloud dimension; however, we leave this population not included because we do not have sufficient knowledge about either the nature of the cloud that holds H$_3^+$ in the (3, 3) level or the mechanism that populates the metastable states, which we discuss separately in the next subsection.  The cloud dimension of the +50~km~s$^{-1}$ wing might not be incompatible with the ``+50~km~s$^{-1}$ clouds'' interpretation in terms of the physical parameters presented by \\citet{gue83}. They estimate the hydrogen column density of M-0.02-0.07, the closest ``+50~km~s$^{-1}$ cloud'' to GC~IRS~3, to be $N$(H$_2$) = 7 $\\times$ 10$^{24}$~cm$^{-2}$, which gives $n$(H$_2$) = 2 $\\times$ 10$^{6}$ cm$^{-3}$ if a cloud dimension of 1~pc is assumed. The angular diameter of M-0.02-0.07 is 2\\arcmin~ in \\citet{gue83}, which is 5~pc at the assumed 8~kpc distance. The apparent smaller pathlength measured by H$_3^+$ could be reconciled if the line of sight toward GC~IRS~3 is off-centered from the cloud core.  On the other hand the physical model of the ``circumnuclear disk'' proposed by \\citet{mar95} gives the radial thickness of the annular disk to be 0.5~pc with a hydrogen density of $n$(H$_2$) $\\simeq$ 10$^{6}$ cm$^{-3}$.  The factor 2--3 difference could be filled by the geometry of GC~IRS~3 and the disk.    \\subsection{H$_3^+$ in Metastable State}  The tentative detection of the $R$(3, 3)$^l$ transition at 3.534~$\\mu$m for the first time in interstellar space would be a breakthrough if confirmed by more observations; it introduces a new dimension in the studies of interstellar H$_3^+$.  The detection of $R$(3, 3)$^l$ and non-detection of other transitions starting from $J >$ 1 levels, such as $R$(2, 1)$^u$, $R$(2, 2)$^l$, $Q$(2, 1)$^l$, $Q$(3, 0) etc., is reasonable since only the (3, 3) rotational level is metastable; spontaneous emission from this level to lower levels is forbidden by the ortho--para selection rule and the absence of the (2, 0) rotational level by the Pauli principle, as shown in figure \\ref{fig:pan2} (Pan, Oka 1986).  On the other hand, H$_3^+$ in the (2, 1), (2, 2), (3, 0), (3, 1), and (3, 2) levels decays to lower levels with lifetimes of 20.4 d, 27.2 d, 3.8 hr, 7.9 hr and 15.8 hr, respectively (Neale et al. 1996) through centrifugal distortion-induced rotational transitions initially proposed for NH$_3$ \\citep{oka71}. Since detailed \\emph{ab initio} quantum chemical calculations of H$_3^+$ have been extensively carried out, these lifetimes are accurate and reliable.  All ortho-H$_3^+$ which are chemically produced or collisionally pumped into the (4, 3), (3, 0), and (3, 3) levels accumulate in the (3, 3) level until it is collisionally deexcited to lower levels.  Similarly, ortho-H$_3^+$ produced in (5, 3), (5, 0), (6, 6), (6, 3), (7, 6), (7, 3), (7, 0) etc. accumulate in the (6, 6) metastable level. Metastable levels are also expected for para-H$_3^+$, for which (4, 4) and (5, 5) may be excessively populated.  H$_3^+$ in the (4, 4) level can spontaneously decay to the (3, 1) level, but its lifetime is long (11 yr).  H$_3^+$ in metastable levels relaxes to lower levels by collisions with H$_2$. Unlike NH$_3$ in which ortho $\\leftrightarrow$ para transitions are forbidden for collision-induced transitions \\citep{che69}, H$_3^+$ may relax to lower levels of different spin modification. This is because a collision of H$_3^+$ with H$_2$ is a chemical reaction and the scrambling of protons may change ortho-H$_3^+$ into para-H$_3^+$ or vice versa, although some selection rules for nuclear spin still remain (Uy et al. 1997; Cordonnier et al. 2000).  There are two ways that the (3, 3) level could be significantly populated compared with the lowest $J$ = 1 levels. First, it may be simply because of high temperature. The (3, 3) level is higher than the lowest (1, 1) level by 251.23~cm$^{-1}$ (361~K) and the populations of the two levels become equal at $T$ = 234~K. Populations in excited non-metastable levels may still be low even at high temperature because of fast spontaneous emission, if the cloud density is lower than the critical densities of the levels. The critical densities are on the order of 10$^4$ cm$^{-3}$ for (2, 1) and (2, 2) and 10$^6$ cm$^{-3}$ for (3, 0), (3, 1) and (3, 2). Thus our observation of the $R$(3, 3)$^l$ transition and non-detection of other transitions starting from $J$ $>$ 1 levels indicate low density. If the metastable state is positively confirmed to be thermally populated, then the $R$(3, 3)$^l$ transition and other absorption lines from metastable levels will make an excellent probe to efficiently isolate the warm interstellar gas. These absorption lines are unique in that they are sensitive exclusively to the high temperature clouds ($T \\sim$ 200~K) along the line of sight with minimum contamination by other cold clouds.  Second, the excess population at (3, 3) could occur at low temperature with a subtle balance of the rate of collisional deexcitation and destruction of H$_3^+$. If the former rate, $k_{\\rm H} n$(H$_3^+$)$n$(H$_2$), is much faster than the latter, $k_{\\rm X} n$(H$_3^+$)$n$(X), H$_3^+$ at low temperature will all relax to the lowest $J$ = 1 levels during its lifetime. The excess population in (3, 3) could occur only when those rates are comparable, that is, $k_{\\rm H} n$(H$_2$) $\\sim$ $k_{\\rm X} n$(H$_{\\rm X}$). This cannot happen in dense clouds where CO is the main destroyer of H$_3^+$, since $k_{\\rm H}$ $\\sim$ $k_{\\rm CO}$ and $n$(H$_2$) $\\gg$ $n$(CO). For diffuse clouds this may be possible if $n$(H$_2$)/$n$(${\\rm e}$) is not much larger than $k_{\\rm e}$/$k_{\\rm H}$ $\\sim$ 250. For this to happen a high electron density and high $k_{\\rm e}$ are preferred. It is interesting to note that the latter requirement is opposite to what is needed to reconcile the enigma of the observed high column densities of H$_3^+$ in a diffuse interstellar medium.  Although our detections of the $R$(3, 3)$^l$ line look fairly convincing for both GC~IRS~3 and GCS~3-2, we note that both lines are broad and without a sharp velocity component. In addition, the spectral line toward GC~IRS~3 is enigmatic in that the five spectral lines starting from the $J$ = 1 level are not strongly observed at the velocity of the $R$(3, 3)$^l$ line. A simple model calculation shows that it is hard to populate the (3, 3) level alone without populating the (1, 0) and (1, 1) levels. It is intriguing that the line profile of $R$(3, 3)$^l$ is similar to those of the strongly saturated CO line both in its wavelength and widths (figure \\ref{fig:h2co}). Such CO lines indicate dense clouds where H$_3^+$ will be collisionally cooled rapidly to lower levels. We may be sampling low-density warm gas with a long pathlength surrounding such dense cloud.  The high population in the (3, 3) level makes the problem of the unexpectedly high abundance in diffuse clouds even harder to reconcile, since other metastable states may also be significantly populated. The need of a high value of $k_{\\rm e}$ for the second mechanism to populate (3, 3) is also an interesting twist related to the enigma. We will attempt to observe more objects and more spectral lines starting from metastable levels, and to carry out a detailed analysis of this problem."
    },
    "0212/astro-ph0212473_arXiv.txt": {
        "abstract": "{We have spectroscopically monitored the galactic Luminous Blue Variable HD\\,160529 and obtained an extensive high-resolution data set that covers the years 1991 to 2002. During this period, the star evolved from an extended photometric minimum phase towards a new visual maximum. In several observing seasons, we covered up to four months with almost daily spectra. Our spectra typically cover most of the visual spectral range with a high spectral resolution ($\\lambda/\\Delta\\lambda \\approx$ 20\\,000 or more). This allows us to investigate the variability in many lines and on many time scales from days to years. We find a correlation between the photospheric He{\\sc i} lines and the brightness of the star, both on a time scale of months and on a time scale of years. The short-term variations are smaller and do not follow the long-term trend, strongly suggesting different physical mechanisms. Metal lines also show both short-term and long-term variations in strength and also a long-term trend in radial velocity. Most of the line-profile variations can be attributed to changing strengths of lines. Propagating features in the line profiles are rarely observed.  We find that the mass-loss rate of HD\\,160529 is almost independent of temperature, i.e.\\ visual brightness. ",
        "introduction": "A-type hypergiants and Luminous Blue Variables (LBVs), which often show spectra similar to A-hypergiants, are the visually brightest stars in galaxies.  The early A-hypergiant \\object{HD\\,160529}, which has long been known to exhibit the spectroscopic signatures of extreme luminosity \\citep{merrill:1943}, has therefore long been considered as one of the visually brightest stars of the Galaxy, comparable to the brightest stars in external galaxies.  The variability of HD\\,160529 was already established 30\\,--\\,40\\,years ago through studies of its variable light and its spectrum at different epochs, which indicated changes in its spectral type.  The spectral variability of HD\\,160529 was studied extensively by \\citet{wolf:1974}, who found variations of the Balmer line profiles, radial velocity variations with an amplitude of about 40\\,km s$^{-1}$\\ and line splitting of some metal lines. Intensity variations of absorption lines of the order of 20\\,\\% and considerable photometric variations were also found. The light variability was studied and described in detail by {\\bf  \\citet{sterken:1977} and \\citet{sterken:1991}.}  HD\\,160529 was classified as a new galactic LBV by \\citet{sterken:1991} due to its brightness decrease of 0.5\\,mag from 1983 to 1991 and an apparent spectral type change from A\\,9 to B\\,8. Apart from the long-term changes, the authors found pulsation-like variations in their photometric data with a quasi-period of 57 days and peak-to-peak amplitudes of 0.1\\,mag in $b$ and $y$, while previous analyses suggested a possible 101.3 day period \\citep{sterken:1981}. More recent photometry has been published by \\citet{manfroid:1991,manfroid:1995} and \\citet{sterken:1993,sterken:1995}.  From comparison with \\object{HD\\,269662} (=R\\,110), a close photometric and spectroscopic counterpart of HD\\,160529 in the Large Magellanic Cloud, \\citet{sterken:1991} derived an absolute visual magnitude of $M_\\mathrm{V} = -8.9$ and a distance of 2.5\\,kpc. They estimated the stellar parameters characterizing the phase of maximum visual brightness to be $T_\\mathrm{eff} \\approx 8\\,000$\\,K, $\\log g \\approx\\,0.55$, $R_\\ast \\approx 330$\\,R$_\\odot$ and $M_\\ast \\approx 13$\\,M$_\\odot$.  The derived mass is in good agreement with the mass determined from evolutionary tracks \\citep{lamers:1998}.  \\citet{sterken:1991} concluded that HD\\,160529 is located near the lower luminosity limit of the LBV instability strip and possibly in an evolutionary phase after the Red Supergiant state, i.e.\\ in a post-RSG evolutionary phase.  HD\\,160529 is of much lower luminosity than more typical LBVs such as \\object{AG\\,Car} \\citep{stahl:2001b} and shows much smaller variations of about 0.5 mag in the visual.  According to \\citet{vangenderen:2001}, it is at the low amplitude limit of the ``strong-active'' S\\,Dor variables. The empirical relation between luminosity and amplitude of LBVs \\citep{wolf:1989} is in agreement with this.  According to the period-luminosity relation of \\citet{stothers:1995} a ``period'' of well above 20 yr would be expected. This is in rough agreement with the light curve published by \\citet{sterken:1991}.  HD\\,160529 can be considered an intermediate case between normal supergiants and LBVs. The closest counterpart may be the LBV \\object{HD\\,6884} (=R\\,40) in the SMC \\citep{szeifert:1993} or the LBV HD\\,269662 (=R\\,110) in the LMC \\citep{stahl:1990}. The variability of HD\\,160529 is therefore of interest both in connection with typical LBVs such as AG\\,Car \\citep{stahl:2001b} and with normal late B -- early A supergiants \\citep{kaufer:1996a,kaufer:1996b}.  Detailed spectroscopic studies of LBVs, which cover both the short and and long time scales, are still very rare and only available for AG\\,Car \\citep{stahl:2001b}. Studies of more objects with different physical parameters are needed in order to distinguish the general behaviour of LBVs from peculiarities of single objects. ",
        "conclusions": "For the interpretation of the short-term variability of HD\\,160529, the relevant time scales are the rotational time scale $P_\\mathrm{rot}$, the stellar-wind time scale $P_\\mathrm{wind}$ and the pulsation time scale $P_\\mathrm{puls}$.  The time scales are defined by the following relations:  \\begin{eqnarray} 2 \\pi R_\\ast / v_\\mathrm{break}  < P_\\mathrm{rot}  <  2 \\pi R_\\ast / v_\\mathrm{rot} \\label{eq:prot} \\\\ P_\\mathrm{wind}       =   R_\\ast/v_\\infty \\label{eq:pwind} \\\\ \\log P_\\mathrm{puls}  =   -0.275 M_\\mathrm{Bol} - 3.918 \\log T_\\mathrm{eff} + 14.543  \\label{eq:ppuls} \\end{eqnarray}  In Equation~\\ref{eq:prot}, the upper limit to $P_\\mathrm{rot}$ is estimated from $v \\sin i$ assuming $\\sin i = 1$. This value can also be slightly larger because the derived $v \\sin i$ value is an upper limit.  {\\bf Note that it is expected that the rotational velocity is a function of phase. Because of the complex line profiles it could not be determined in light minimum and is therefore set constant here.} Because of the acceleration of the wind to its terminal velocity in the extended stellar wind, the typical expansion time scale can be significantly longer than given by Equation~\\ref{eq:pwind}.  The pulsation period for the fundamental radial mode given by Equation~\\ref{eq:ppuls} is the empirical fitting formula given by \\citet{lovy:1984}.  We used the following parameters from \\citet{sterken:1991}: $M_\\mathrm{bol} = -8.9$, $M_\\ast = 13\\,\\mathrm{M}_\\odot$.  In addition, we used $v_\\infty$ = 200 km s$^{-1}$ and $v_\\mathrm{rot}$ = 45 km s$^{-1}$ for both maximum and minimum phase. For the stellar radius and temperature we adopted 150 R$_\\odot$/12\\,000~K and 330 R$_\\odot$/8\\,000~K for minimum and maximum phase, respectively.  With these parameters, we obtain the time scale estimates given in Table~\\ref{periodtab}.  \\begin{table} \\caption[]{Summary of time scale estimates (in days) for the minimum and maximum phase of HD\\,160529.} \\begin{center} \\begin{tabular}{lll} & min. (hot) & max. (cool) \\\\ \\hline $P_\\mathrm{rot}$  & 59 \\ldots 169 & 193 \\ldots 371  \\\\ $P_\\mathrm{wind}$ & $>$6          & $>$13           \\\\ $P_\\mathrm{puls}$ & 10            & 49              \\\\ \\end{tabular} \\label{periodtab} \\end{center} \\end{table}  Within the uncertainties, the time scales $P_\\mathrm{wind}$ and $P_{\\rm puls}$ are compatible with the typical variation time scales of 50--100 days, while the rotational time scale is significantly longer.  However, {\\bf since rotational modulation can also lead to variations with an integer fraction of the rotational period, the time scale of the variations cannot be used to distinguish the possible mechanisms of the short-term variations in HD\\,160529.} The simultaneous variation of the optical brightness and the He{\\sc i}$\\lambda$5876 line, however, is most easily explained by pulsations. Pulsations have also been suggested as a cause for the photometric micro-variations of LBVs \\citep[cf.~e.g.][]{lamers:1998}.  The correlation of the equivalent width variations with visual brightness on short time scales is in marked contrast with the behaviour on long time scales. This strongly suggests that the physical mechanism for the variations on short time scales is different from the LBV-type variations.  It is rather surprising, that so far no strong dependence of the observed time scale on the visual brightness -- and therefore the radius -- has been detected. All plausible variability mechanisms would predict such a dependence. It should be mentioned, however, that \\citet{sterken:1991} derived a period of 57 days in minimum while the derived value in maximum was 101 days -- about a factor of two longer. It is not clear, however, if the difference is due to a real period change.  Most LBVs seem to increase their mass-loss rate in maximum.  However, it is not clear that this is a general property of LBVs \\citep{leitherer:1997}. \\citet{vink:2002} have shown that the mass-loss rate of LBVs versus temperature carries important diagnostic information. The mass-loss rate of \\object{AG\\,Car} versus phase has been studied by \\citet{stahl:2001b}. \\citet{vink:2002} compared these results with their predictions for line-driven stellar winds and found a good agreement. They also predict that the mass-loss rate of LBVs is a complicated function of temperature, surface gravity and abundances. In particular, the mass-loss rate can even decrease with decreasing temperature at temperatures below about 15\\,000 K, in particular at low masses. This predicted behaviour nicely fits the observed behaviour of HD\\,160529, which is an LBV with very low temperature and mass.  Although the mass-loss rates of LBVs can be predicted with some precision, the details of the mass-loss process are still not understood. In particular the complicated line-profile variations and the line splitting cannot be explained by radially expanding features. Rotational modulation may be the cause for these variations."
    },
    "0212/astro-ph0212184_arXiv.txt": {
        "abstract": "We use {\\it Hubble Space Telescope} ({\\it HST}) near-infrared imaging to explore the shapes of the surface brightness profiles of bulges of S0-Sbc galaxies at high resolution.  Modeling extends to the outer bulge via bulge-disk decompositions of combined {\\it HST} - ground based profiles.  Compact, central unresolved components similar to those reported by others are found in $\\sim$84\\% of the sample.  We also detect a moderate frequency ($\\sim$34\\%) of nuclear components with exponential profiles which may be disks or bars.  Adopting the S\\'ersic \\rn\\ functional form for the bulge, none of the bulges have an \\r14\\ behaviour; derived S\\'ersic shape-indices are $<n> = 1.7 \\pm 0.7$.  For the same sample, fits to NIR ground-based profiles yield S\\'ersic indices up to $n = 4-6$.  The high-$n$ of ground-based profiles are a result of nuclear point sources blending with the bulge extended light due to seeing.  The low S\\'ersic indices are not expected from merger violent relaxation, and argue against significant merger growth for most bulges. ",
        "introduction": "\\label{Sec:Introduction}  Whether galaxy mergers (e.g.\\ Kauffmann et al.\\ 1996) or instead disk instabilities (Pfenniger \\& Norman 1990) may be viewed as the dominant process of formation and growth of the central bulges in spiral and S0 galaxies depends to some extent on the shape of the surface brightness profiles of bulges.  Simulations of mergers of similar-size galaxies yield spheroidal systems following the \\r14\\ law (e.g.\\ Barnes 1988). The accretion of dense satellites onto bulge-disk galaxies also leads to \\r14\\ bulges (Aguerri, Balcells \\& Peletier 2001, herafter ABP01). In contrast, disk instabilities lead to bulges with approximately exponential profiles (Combes et al.\\ 1990).  Over the past decade several studies have stressed that bulges of late- to intermediate-type spiral galaxies are best fit with profiles of exponential to $r^{1/2}$ shapes (Andredakis \\& Sanders 1994; de Jong 1996; Courteau, de Jong \\& Broeils 1996; Seigar \\& James 1998; Carollo, Stiavelli \\& Mach 1998; MacArthur et al.\\ 2002).  However, many bulge-disk decompositions are still performed using \\r14\\ bulges only, especially at high-$z$, e.g. Simard et al.\\ (2002); the \\r14\\ law is often cited for massive or early-type bulges: Carollo and collaborators (e.g.\\ Seigar et al.\\ 2002 and references therein) routinely classify bulges into \\r14\\ and exponential types. Do massive bulges follow the \\r14\\ model? Do bulges come in two flavors, perhaps pointing at massive bulges growing by early mergers and less-massive bulges growing by disk instabilities?  Fitting bulge profiles with the S\\'ersic model $\\mu(r) \\sim r^{1/n}$ provides an avenue for measuring the true shapes and concentrations of the light distribution of bulges.  Various studies that include early-type bulges (Andredakis, Peletier \\& Balcells 1995, hereafter APB95; Khosroshahi et al.\\ 2000; Graham 2001; M\\\"ollenhoff \\& Heidt 2001) find a continuous distribution of S\\'ersic shape indices $n$ that scales with bulge luminosity from $n<1$ to $n>4$.  At the high-$n$ end, these results depend strongly on the innermost regions of the profiles.  Inner disks, compact sources and star formation, present in galaxy nuclei (e.g.\\ Phillips et al.\\ 1996; Carollo et al.\\ 2002; Rest et al.\\ 2001; Pizzela et al.\\ 2002), contribute light that cannot be disentangled from the spheroid light at ground-based resolutions (Ravindranath et al.\\ 2001).  Extra central light results in higher-$n$ profile fits.  Nineteen galaxies from the Balcells \\& Peletier (1994, BP94) sample were imaged with {\\it HST}/NICMOS at F160W (Peletier et al.\\ 1999, Paper I).  We have performed bulge-disk decompositions, and analyze profiles, global parameter relations and nuclear cusp slopes in a companion paper (Balcells et al.\\ 2002, Paper III).  In this Letter, we focus on the shape indices obtained from S\\'ersic fits to the bulge profiles.  The high resolution of the {\\it HST} data allows one to disentangle nuclear disks/bars and compact sources from the extended spheroid light.  Modeling these components, we study S\\'ersic indices for the bulge profiles and compare them to those obtained from ground-based profiles.  We characterize the magnitude distribution of point sources found in the galaxy centers and discuss the relevance of low-$n$ profiles for bulge formation models.  Preliminary results on nuclear cusps are presented in Balcells et al.\\ (2001).  A key ingredient of the analysis of bulge profiles, overlooked in the {\\it HST}-based studies of bulges cited above, is the modeling of the extended bulge and disk light distributions when fitting analytic functions to the bulge light.  Here, composite profiles linking {\\it HST} profiles with ground-based $K$-band profiles are used so that bulge-disk decompositions can be performed.  We assume a Hubble constant of $H_{0} = 75$ \\kms\\,Mpc$^{-1}$. ",
        "conclusions": "\\label{Sec:Discussion}  Measuring the correct value of the S\\'ersic index $n$ for bulges is important for several reasons.  Overestimating $n$ leads to an overestimate of the bulge $r_e$ and luminosity.  Trujillo et al.\\ (2001) show that fitting an $r^{1/2}$ profile with a \\r14\\ function yields an overestimate of $r_e$ by a factor of 3-5 and of the bulge luminosity by a factor of about 2.  The luminosity ratio B/D is further overestimated as more of the light in the bulge-disk transition region is assigned to the bulge.  The converse occurs when $n$ is underestimated, eg.\\ when exponential fits are forced onto $n>1$ S\\'ersic profiles (Graham 2001).  Thus, small luminosity components (typically $L_{K,PS}/L_{K,Bul} \\approx 10^{-2.8}$, eqn.\\ 1) may significantly alter the derived global parameters of bulges.  Test fits excluding inner profile regions suggest that, in the absence of {\\it HST} imaging, the effects of central components on the bulge shape index $n$ are greatly diminished by excluding the central FWHM of the profile, as done by Stiavelli et al.\\ (2001) for dwarf ellipticals.  Point sources should present less of a problem for late-type spirals given the low S\\'ersic indices, low bulge luminosities and hence low PS luminosities expected from Figure~\\ref{Fig:PtSrc}c.  The nuclear PSs, amounting to a fraction of order $10^{-3}$ of the bulge luminosities, are easy to accomodate within either a merger origin or a secular evolution origin for bulges.  The observed scaling of PS and bulge luminosities need not indicate an internal origin for the PS as nuclear gas deposition during a merger may scale with the masses of the merging objects.  If bulges are evolved bars, the problem of growing a nuclear star cluster is similar to that of feeding an active galactic nucleus (Maciejewski et al.\\ 2002).  Given the persistent description of bulges as \\r14\\ spheroids in galaxy formation models, it is worth pointing out that none of the bulges follow the \\r14\\ law.  APB95, de Jong (1996), Graham (2001) and MacArthur et al.\\ (2002) have shown that late- to intermediate-type bulges have $r^{1/n}$ profiles in the range $0.5 \\lesssim n \\lesssim 2.0$.  Our work establishes a similar range, $0.5\\lesssim n \\lesssim 3.0$, for intermediate- to early-type bulges i.e. those which are still generally described with the \\r14\\ model. Extrapolating from the rising envelope of the $n$ -- $M_K$ relation in Figure~\\ref{Fig:NvsTvsBDvsMKdeJ}d, only bulges with $M_K \\leq -26$ may reach $n=4$.  The \\r14\\ law is interpreted as the consequence of violent relaxation during a clumpy collapse (van Albada 1982), and is obtained after mergers of similar-size disk galaxies (Barnes 1988).  Simulations of collisionless satellite accretion onto bulge-disk galaxies further show that bulges evolve smoothly from exponential profiles to \\r14\\ profiles when doubling their mass due to accretion events (ABP01).  At face value then, our distribution of S\\'ersic shape indices $n$ (Figure~\\ref{Fig:NvsTvsBDvsMKdeJ}c,d) suggests that the vast majority of bulges did not acquire their present structural properties by growing through collisionless mergers.  Kinematic data might show whether, besides a low S\\'ersic index $n$, these bulges show other properties typical of disks (Kormendy 1993; Falc\\'on-Barroso et al.\\  2002).   Have deviations from the \\r14\\ law in $N$-body models gone undetected?  S\\'ersic fitting is only starting to be applied in $N$-body studies.  Hjorth \\& Madsen (1995) show that violent relaxation in a finite volume, combined with subsequent escape of positive-energy stars, leads to a distribution function with deviations from the \\r14\\ law; from their figures, a shallow central potential might possibly fit $n<4$ S\\'ersic profiles. Accretion of low-mass satellites might heat the bulge only moderately and deposite the compact satellite nucleus at the center of the remnant.  This could lead to a $n\\sim 1-2$ bulge with a central brightness spike not unlike what is observed.  But very long merger times and slow growth rate for low-mass satellites are a problem for this hypothesis.  Recent cosmological SPH simulations of galaxy formation (Scannapieco \\& Tissera 2002) find $n\\approx 1$ for bulges resulting from secular evolution, while merging rises $n$ to about $n=4$, again pointing at low-$n$ profiles being related to non-merger formation histories."
    },
    "0212/astro-ph0212071_arXiv.txt": {
        "abstract": "We present the results of a large grid of synthetic spectra and compare them to early spectroscopic observations of SN~1993W. This supernova was discovered close to its explosion date and at a recession velocity of 5400~\\kmps\\ is located in the Hubble flow. We focus here on two early spectra that were obtained approximately 5 and 9 days after explosion.  We parameterize the outer supernova envelope as a power-law density profile in homologous expansion.  In order to extract information on the value of the parameters a large number of models was required. We show that very early spectra combined with detailed models can provide constraints on the value of the power law index, the ratio of hydrogen to helium in the surface of the progenitor, the progenitor metallicity and the amount of radioactive nickel mixed into the outer envelope of the supernova. The spectral fits reproduce the observed spectra exceedingly well. The spectral results combined with the early photometry predict that the explosion date was $4.7 \\pm 0.7$~days before the first spectrum was obtained. The ability to obtain the metallicity from early spectra make SN~IIP attractive probes of chemical evolution in the universe and by showing that we have the ability to pin down the parameters of the progenitor and mixing during the supernova explosion, it is likely to make SN~IIP useful cosmological distance indicators which are at the same time complementary to SNe~Ia. ",
        "introduction": "Supernova 1987A (SN~1987A) confirmed our basic theoretical understanding of a Type II supernova (SN~II) as the core collapse of a massive star, which leaves behind a compact object (neutron star or black hole) and expels the outer mantle and envelope into the interstellar medium \\citep[see][and references therein]{sn87arev89}. While the explosion mechanism itself remains a subject of active research, the propagation of the shock wave through the mantle and envelope is reasonably well understood. Theoretical models for the light curve do a good job of reproducing the observations \\citep{blinn87a99a,blinn87a00} and detailed radiation transport calculations verify that these models reproduce the observed spectra from the UV to the IR \\citep{mitchetal87a01,mitchetal87a02}. Unlike Type Ia supernovae (SNe~Ia) where the actual compositions as a function of velocity are an important subject of current research \\citep{fisher91T99}, the compositions of the envelopes of red or blue supergiant stars is primarily hydrogen and helium and thus we show that the primordial compositions, and the degree of mixing of hydrogen and \\nni, can be determined to high precision by the detailed spectral modeling of observed SN~II spectra.  SNe~II have a very large spread in their intrinsic brightness, from the very dim SN~1987A, to the exceedingly bright SNe~1979C and 1997cy. The observed spread in intrinsic luminosity is greater than a factor of 500. This is not surprising given the fact that the progenitors span a wide range of initial stellar masses, possible binary membership, and prior star formation histories. Clearly, SNe~II do not meet the astronomers requirement of being a \\emph{standard candle}, however we believe that the fact that our models can determine the stellar compositions, degree of mixing, and kinetic energy of the explosion shows that their atmospheres can be well understood. Although we do not present distances in this paper, we believe that these results increase the attractiveness of SNe~II as cosmological probes and in the future SNe~II will become complementary with SNe~Ia as distance indicators. Both the SEAM \\citep[see][and references therein]{b93j4,b94i1,mitchetal87a02} and the EPM \\citep[see][]{baadeepm,skeetal94,hamuyepm01,leonard99em02} methods for determining distances to SNe~II (and other supernova types for SEAM) depend on the ability to model the spectral energy distribution (SED) of the supernova atmosphere accurately. Clearly, quantities like the composition and degree of mixing play a role in the output SED and thus the work presented here helps to place both these methods on firmer footing.  While the average SN~II is several times dimmer than a SN~Ia, current ground-based searches and proposed future space-based searches for supernovae will easily detect these objects at cosmologically interesting distances. In the very deep high-redshift supernova searches by the Hi-Z Team \\citep{riess_scoop98} and the Supernova Cosmology Project (SCP) \\citep{perletal99} several SNe~II have been discovered even though the searches are geared to finding SNe~Ia. The SCP has positively identified via spectroscopy 5 SNe~II out to a redshift $z \\approx 0.45$. \\citet{hdf_gnp99} detected SN~1997fg through differencing two epochs of Hubble Deep Field North $I$-band photometric images. While the supernova was not spectroscopically classified, the combination of the host galaxy type (an irregular), its redshift $z = 0.952$, and the supernova's apparent brightness makes it a strong candidate for a SN~II. A. Clocchiatti \\etal\\ (IAUC 7549) found SN~2000fp at $z=0.30$, G. Altavilla \\etal\\ (IAUC 7762) found SN~2001gg at $z=0.61$, SN~2001gh at $z\\approx 0.16$, and SN~2001gj at $z=0.27$, and G. Altavilla \\etal\\ found SN~2002co at $z=0.318$. The proposed \\emph{SNAP} satellite, a wide-field optical imaging telescope, (see {\\tt http://snap.lbl.gov}) will easily be able to find, follow, and spectroscopically identify SNe~II beyond redshifts of $z \\approx 1$. \\emph{NGST}, a large, small field IR telescope, will be able to detect SNe~II to the initial star formation period in the universe, $z \\ga 5$.  The explosion mechanism is now thought to result in a non-spherically symmetric shock-wave \\citep{lbw70,khok_jet99} and this has been confirmed by the detection of significant polarization in the spectra of SNe~II \\citep{wang_pol96,leonmd00}. Nevertheless, the large expansion of the ejecta before the supernova becomes visible leads to significant ``sphericalization'' \\citep{chev84}, and asphericity effects may only be important at late times, or in SNe~II that are strong circumstellar interacters (SNe~IIn). These supernovae are clearly distinguishable observationally from ``normal'' Type II supernovae.  Calculations \\citep{miralda97,DF99,livyoung00} indicate that SNe~II may in fact be the first stellar objects visible in the universe and hence they serve as important probes for the star formation rate, the rate of chemical evolution by measuring their primordial abundances, and directly for the cosmological parameters.  We show that the primordial abundances of the progenitor star can be reliably determined by detailed synthetic spectral modeling. ",
        "conclusions": "We have shown that synthetic spectra of SNe~IIP can be fit to high accuracy, and that the primordial metallicity, degree of mixing of H and He, as well as the amount of mixing of \\nni\\ can also be determined by early spectra. \\citet{bsn99em00} showed that the total amount of extinction due to dust either in the parent galaxy or in our own galaxy can also be estimated with UV+optical spectra at very early times. While the reduced calcium abundance that was needed to reduce the strength of the calcium lines is intriguing, we do not as yet believe that our results require a reduced value of Ca/Fe.  In a future paper we will explore a grid of models and determine a distance to SN~1993W, along with an estimate of the systematic error. We have shown that primordial metallicity, H/He mixing, and \\nni mixing are derivable from the detailed analysis of early spectra of SNe IIP. The analysis of abundances of individual elements is in reach, but will require a larger grid of models intercompared with a set of well studied SNe IIP. The data for this will be obtained as part of searches for nearby SNe~Ia and will just require followup spectrophotometry. We believe that we have demonstrated that SNe~IIP atmospheres can be well understood theoretically and hence are likely to make excellent independent cosmological probes, both as distance indicators and as probes of cosmic nucleosynthesis."
    },
    "0212/astro-ph0212247_arXiv.txt": {
        "abstract": "We present 86.8 GHz SiO J=2-1 observations of the molecular outflow from the protostellar source NGC 1333 IRAS 2A. Silicon monoxide is a sensitive tracer of shocks in molecular gas, and the emission clearly traces terminal bowshocks at the ends of the outflow lobes, as well as emission knots and a source coincident with the protostar. Our three-field mosaic from the BIMA interferometer covers the entire outflow length, including both lobes.  The most prominent emission feature, the eastern bowshock, shows SiO at the leading edge, clearly displaced from the maximum  seen in CO. This is in contrast to recent observations by Chandler \\& Richer (2001) of  the HH 211 jet, where SiO traced primarily internal working surfaces and was less extended than the CO.  Our results are consistent with prompt entrainment at the leading edge of the IRAS 2A bowshocks. ",
        "introduction": " ",
        "conclusions": "The class 0 source IRAS 2A drives an unusual, highly collimated jet-like outflow first studied in CO (Sandell {\\it et al}.~1994). Such a young, symmetric, and focused outflow provides an excellent opportunity to observe the expected bow shock of a jet interacting with the molecular ambient gas.  We used the BIMA interferometer to map this outflow in SiO in order to trace shocks within the outflow. Our initial three-field, C-array  (10''$\\times 7''$ beam) mosaic covers the positions of IRAS 2A and the east and west bow shocks. In addition to this outflow of roughly E-W orientation, a larger, less collimated bipolar outflow originates from nearly the same position (Liseau {\\it et al.}~1988), and is possibly associated with the nearby source IRAS 2C (Knee \\& Sandell 2000; Sandell \\& Knee 2001). Because the position of IRAS 2C agrees poorly with the symmetry axes of the flows seen in CO, the possibility exists that the two perpendicular outflows seen in CO might actually represent a single, much wider, limb-brightened  bipolar outflow (see Figure~1). \\begin{figure}[h] \\begin{center} \\centerline{\\includegraphics[width=30pc,angle=270]{wilkin_fig1.eps}} \\caption{Mosaic of the IRAS 2A region. Thin contors show 12CO J=2-1 taken from the BIMA observations of Engargiola \\& Plambeck 1999. The greyscale indicates blue-shifted CO. Thick contors show SiO J=2-1 emission. Crosses of decreasing size give the positions  of IRAS 2A, 2B, and 2C.} \\label{fig:map} \\end{center} \\end{figure} Our recent observations with BIMA in the D-array (13.5'' beam) detect SiO in the E-W outflow but not the N-S, strongly suggesting that these are two separate outflows. Nonetheless, it still appears likely that the driving source of the N-S flow remains to be discovered, and further searches for binarity of IRAS 2A are suggested.  In Figures~2 and 3 we present the higher resolution wideband and channel maps of SiO in the E-W outflow. The much brighter, redshifted, eastern bow shock shows a wider range of velocities, and is associated with a dense clump seen in CS by Sandell {\\it et al.}, and whose V-shaped geometry resolved by Blake (1996) suggests a bow shock driven into a steep density gradient. Presumably the eastern bow shock is much brighter because the driving jet is interacting with this dense clump. The actual emission peaks in both the east and west lobes are unresolved in our maps, and higher resolution observations are needed to confirm the  expected bow shock geometry at the ends of the outflow lobes. \\begin{figure}[h] \\begin{center} \\centerline{\\includegraphics[width=30pc,angle=270]{wilkin_fig2.eps}} \\caption{3-Field mosaic of IRAS 2A outflow in SiO j=2-1. In this map, positions are relative to the bright, eastern emission peak. The submm point source locations are shown in Figures~1 and 3.} \\label{fig:map} \\end{center} \\end{figure} \\begin{figure}[h] \\begin{center} \\centerline{\\includegraphics[width=30pc,angle=270]{wilkin_fig3.eps}} \\caption{Channel maps at $5\\,{\\rm km}/{\\rm s}$ intervals. The large, medium and small crosses indicate the positions of IRAS 2A, 2B and 2C, respectively.} \\label{fig:channels} \\end{center} \\end{figure}"
    },
    "0212/astro-ph0212137_arXiv.txt": {
        "abstract": "A set of simulations concerning the influence of internal energy on the stability of relativistic jets is presented. Results show that perturbations saturate when the amplitude of the velocity perturbation approaches the speed of light limit. Also, contrary to what predicted by linear stability theory, jets with higher specific internal energy appear to be more stable. ",
        "introduction": "The production of collimated relativistic outflows is common among extragalactic radio sources. The question of why some of these sources produce jets that propagate up to hundreds of kpcs along nine decades in distance scale (as, e.g., Cyg~A) whereas others (e.g., 3C31) decollimate and flare after few kpc, arises naturally. To give an answer, the analysis of the nonlinear stability of relativistic flows against the growth of Kelvin-Helmholtz (KH) perturbations emerges as a powerful tool.  Focusing on the relativistic regime, the linear growth of KH perturbations both in the vortex sheet approximation (see Birkinshaw 1991a for a review), and sheared flows (Birkinshaw  1991b, Hanasz \\& Sol 1996) has been extensively analyzed. However, the nonlinear regime remains practically unexplored with only a couple of papers (i.e., Rosen et al. 1999, Hardee et al. 2001) covering partial aspects of the problem. In the present work, we concentrate on the influence of the specific internal energy in the long term stability of relativistic flows. The reason for this relies on two points: i) it is a genuine relativistic effect and no numerical study has been yet performed for it (Rosen et al. 1999 have reported preliminary results), and ii) a remarkable stability of extremely {\\it hot} jets has been found in simulations (see, e.g., Mart\\'{\\i} et al. 1997), in contradiction with predictions of linear stability analysis. ",
        "conclusions": ""
    },
    "0212/astro-ph0212301_arXiv.txt": {
        "abstract": "The expansion of the universe causes spacetime curvature, distinguishing between distances measured along and transverse to the line of sight. The ratio of these distances, e.g.~the cosmic shear distortion of a sphere defined by observations of large scale structure as suggested by Alcock \\& Paczy{\\'n}ski, provides a method for exploring the expansion as a function of redshift.  The theoretical sensitivity to cosmological parameters, including the dark energy equation of state, is presented. Remarkably, sensitivity to the time variation of the dark energy equation of state is best achieved by observations at redshifts $z\\la 1$.  While systematic errors greatly degrade the theoretical sensitivity, this probe may still offer useful parameter estimation, especially in complementarity with a distance measure like the Type Ia supernova method implemented by SNAP.  Possible future observations of the Alcock-Paczy{\\'n}ski distortion by the KAOS project on a 8 meter ground based telescope are considered. ",
        "introduction": "\\label{sec.intro}  We now have strong evidence that the expansion of the universe is accelerating, from the original method of Type Ia supernova distance-redshift measurements \\cite{perl99,riess98} and concordant observations of the cosmic microwave background (CMB) power spectrum and of large scale structure \\cite{bond,perc}. Understanding the nature of the dark energy responsible for the acceleration will have profound implications for cosmology, high energy physics, and fundamental physics. Mapping the expansion history of the universe offers a way to gain insights into the dark energy and the fate of the universe, for example by characterizing the equation of state behavior which is directly related to properties of the scalar field potential.  Distance measures, notably the supernova method, have proved useful at constraining the energy density and equation of state of the dark energy, with great improvements expected in the next decade.  But these involve an integration over the expansion rate behavior $H(z)$, which itself involves a redshift integral over the equation of state $w(z)$.  We can ask whether we can devise a more direct probe of the acceleration.  In fact one such, the redshift drift test, was proposed by Sandage \\cite{san61} in 1961 and developed further by Linder \\cite{lin97}.  The redshift of a source is a central astrophysical observable.  It is directly related to the change in time intervals due to the cosmic expansion between a photon's emission and observation, $z=dt_o/dt_e-1=a_o/a_e-1$, where $a(t)$ is the scale factor of the universe.  But obviously one could consider a second derivative term, a time dependence of the redshift itself as the universe ages: \\beqa {dz\\over dt_o}&=&{d\\over dt_o}\\left[{a(t_o)\\over a(t_e)}\\right]=[\\dot a(t_o)- \\dot a(t_e)]/a(t_e) \\\\ &=&H_0(1+z)-H(z). \\label{dzt}\\eeqa  This provides a direct measure of acceleration, being effectively a second time derivative, as can be seen from the low redshift limit: $\\Delta z/z\\approx -q_0H_0\\Delta t$, where $q_0$ is the present deceleration parameter. However since the astronomical observing time is much smaller than the Hubble time, $\\Delta t\\ll H_0^{-1}$, this is not a practical probe.  For example in the most optimistic case of observing over a period of 10 years a hypothetical spectral line emitted at the CMB last scattering surface at $z=10^3$, one requires a redshift measurement of a part in $10^5$ to distinguish cosmological models.  For emission line objects at $z=5$, this becomes a part in $10^8$.  Moreover, just as peculiar velocities affect redshift measurements at the level of $v/c$, so do peculiar accelerations, i.e.~local gravitational potentials $\\Phi$ from inhomogeneously distributed matter, affect the redshift drift measurements at a level $\\Phi/c^2\\approx10^{-5}$. This latter effect even ruins the generalization of the redshift drift called the cosmic pulsar test, where timing is improved by measuring a large number $N$ of wavelengths or pulses \\cite{lin97}.  But if we are stymied in measuring the acceleration directly, at least we can hope to measure the first derivative of the expansion, $H(z)$.  One of the ways to do this is the cosmic shear, or {\\alp} test.  Section \\ref{sec.form} sets up the formalism while Section \\ref{sec.sens} applies Fisher matrix analysis to investigate the theoretical sensitivity of this method for estimating the cosmological parameters.  In Section \\ref{sec.sys} we introduce observational sanity in the form of systematic uncertainties and discuss the proposed KAOS project as a means of carrying out this test.  We summarize our conclusions and plans for future work in Section \\ref{sec.concl}. ",
        "conclusions": "\\label{sec.concl}  The cosmic shear test looks extremely promising theoretically for the determination of cosmological parameters.  This is evident from its tomographic dependence on the expansion rate of the universe, shown by the appearance of the Hubble parameter $H(z)$ by itself.  It has further interesting properties in the redshift evolution of its parameter degeneracies and complementarity with other probes.  However, systematic uncertainties, most probably involving peculiar velocities and distortions related to the local environment rather than cosmic expansion, put severe limits on the probe's ability to fulfill its potential.  Unless systematic uncertainties can be brought under the 2\\% level, the cosmic shear test applied at low or high redshift does not appear to offer significantly complementary or generally comparable limits to the supernova distance method.  Considering a linear systematic of $0.02z$, a matter density prior $\\sigma(\\Omega_m)=0.01$, or the addition of Planck CMB data does not greatly affect these conclusions.  The best hope for applying this method does, however, lie at redshifts $z<1$, which is observationally feasible. Therefore careful study of the systematic uncertainties might yield some regime in which the next generation of wide, deep redshift surveys can bring this method of determining cosmological parameters to fruition."
    },
    "0212/astro-ph0212393_arXiv.txt": {
        "abstract": "A new analysis of Keck/HIRES observations of the broad absorption line QSO \\apm\\ indicates that a number of intervening \\cfour\\ absorbers give rise to absorption lines for which the observed optical depths for 1548, 1550~\\AA\\ doublet components are not in the expected $2:1$ ratio. To compensate for the effect, a local adjustment of the zero-level is required. We model this effect as coverage of one line of sight to this gravitationally lensed QSO and perform a set of simulations to select a sample of lines for which our model provides an explanation for the effect. We use lines in this sample to obtain estimates for minimum \\cfour\\ absorber sizes from total coverage and the separations of the lines of sight for a range of lens redshifts, $z_{\\rm l}$, and two cosmologies. We also obtain best estimates for overall sizes from a statistical \\lq hit and miss\\rq\\ approach. For $z_{\\rm l}=0.7$ our results set a lower limit to sizes of \\cfour\\ absorbers of $\\sim 0.3 \\, h^{-1}_{72}$ kpc ($\\sim 0.5 \\, h^{-1}_{72}$ kpc) for $\\Omega_M=1, \\Omega_{\\Lambda}=0$ ($\\Omega_M=0.3$, $\\Omega_{\\Lambda}=0.7$), in agreement with other results from similar work but are limited by sample size and the uncertainty in $z_{\\rm l}$. Our method can be used to detect lensed QSOs and to probe absorber sizes when separate spectra cannot be obtained for each line of sight. ",
        "introduction": "The advent of 10-m class telescopes in the last decade has provided us with an unprecedented view of the high-redshift universe. In particular, HIRES, the high-resolution echelle spectrograph on the Keck I telescope in Hawaii, has made it possible to obtain observations with resolution as high as $\\sim 6$~\\kmps. As a result, a number of detailed studies of absorption lines in the spectra of background QSOs have been carried out (for reviews see Rauch 1998; Hamann \\& Ferland 1999).  However, this improvement in {\\em spectral} resolution cannot be easily accompanied by an improvement in {\\em spatial} resolution: 2-dimensional information about the absorbers can normally only be obtained by comparing separate spectra of multiple quasar images, which are due to either gravitationally lensed single QSOs or actual QSO pairs which have a small angular separation on the plane of the sky. Because lensed QSOs have image separations up to a few arcseconds, they are useful for probing scales $\\approxlt 100$~kpc. QSO pairs have image separations of at least a few arcminutes, thus allowing one to probe scales of a few hundred kpc. From coincident absorption in more than one line of sight (LOS), minimum size (coherence length) estimates can be obtained immediately. Estimates of most probable sizes require a statistical \\lq hit and miss\\rq\\ approach (e.g. McGill 1990).  Such studies have led to size estimates for \\lya\\ absorbers ranging from a few kpc (Foltz et al. 1984; McGuill 1990; Bechtold \\& Yee 1995) to a few hundred kpc (Dinshaw et al. 1995; Smette et al. 1995; Dinshaw et al. 1998) to $\\sim 1$~Mpc (Fang et al. 1996; Dinshaw et al. 1997; Young et al. 2001). Such large \\lq sizes\\rq\\ suggest that the typical, low column density ($\\approxlt 10^{14.5}$\\pcmsq) \\lya\\ forest line does not originate in discrete galactic or protogalactic clouds. The emerging picture is consistent with results from hydrodynamical numerical simulations in a Cold Dark Matter (CDM) structure formation scenario according to which the \\lya\\ forest is a manifestation of a space filling photoionized intergalactic medium (IGM). This is the most important baryon reservoir in the Universe at high redshift (see e.g. Efstathiou, Schaye \\& Theuns 2000 for a review).  The \\cfour\\ doublet ($\\lambda \\lambda 1548.20, 1550.78$~\\AA) is the most commonly observed heavy-element absorption signature redward of the \\lya\\ emission line in spectra of high-redshift QSOs.  The oscillator strengths for the 1548.20 and 1550.78~\\AA\\ components are in a $2:1$ ratio, which translates to the same optical depth ratio for optically thin ($\\tau \\ll 1$) lines.  Line parameters are consistent with a photoionized gas (Rauch et al. 1996).  The \\cfour\\ absorbers appear to be highly clustered (Sargent et al. 1988; Petitjean \\& Bergeron 1994; Loh, Quashnock \\& Stein 2001) and have been thought to originate in gas clouds orbiting in galactic haloes (Sargent et al. 1979; Steidel 1993; Petitjean \\& Bergeron 1994; Mo \\& Miralda-Escud{\\'e} 1996). Hydrodynamical simulations also support a picture in which high-redshift \\cfour\\ arises in protogalactic clumps (PGCs), i.e., small, future constituents of large galaxies, forming through gravitational instability in a hierarchical cosmogony (Haehnelt, Steinmetz \\& Rauch 1996; Rauch, Haehnelt \\& Steinmetz 1997).  Size estimates for heavy element absorbers range from a few kpc (Young et al. 1981; Smette et al. 1992; Monier, Turnshek \\& Lupie 1998) to a few tens of kpc (Crotts et al. 1994; Smette et al. 1995; Lopez, Hagen \\& Reimers 2000) or even hundreds of kpc (Shaver \\& Robertson 1983; Petitjean et al. 1998). A recent study using three gravitationally lensed QSOs (Rauch, Sargent \\& Barlow 2001) obtains size estimates on the order of kiloparsecs. It is also apparent that sizes are absorber-type dependent, with individual clouds for low-ionization absorbers having dimensions on the order of a few $\\times 10$~pc as opposed to at least a few $\\times 100$~pc for high-ionization absorbers (Rauch, Sargent \\& Barlow 1999; Rauch et al. 2001; Rigby, Charlton \\& Churchill 2002).  With an $R$-band magnitude of 15.2 (Irwin et al. 1998) \\apm\\ is one of the intrinsically most luminous QSOs known (bolometric luminosity $\\approxlt 10^{15} L_{\\odot}$, Egami et al. 2000; Ibata et al. 1999). From molecular CO emission lines the estimated quasar redshift is $z=3.911$ (Downes et al. 1999). It is thus well suited for QSO absorption line studies. Additionally, the object is gravitationally lensed (Ledoux et al. 1998; Ibata et al. 1999; Egami et al. 2000) but the redshift of the lens is not known. As the object is lensed into three images, it represents the first confirmed odd-image lens system (Lewis et al. 2002).  For this object it is not possible to obtain separate spectra for different LOS with ground-based observations (image separation $\\sim 0.38$~arcsec). However, if only one LOS is covered, there is an excess in the continuum intensity that gets through to the observer.  If the absorbing column is high enough, the observed optical depths of the components of a heavy element doublet are not in the expected ratio, providing a signature for this effect. We use the term Anomalous Doublet Ratio (ADR) for such doublets. In such a case, in order to obtain a satisfactory fit, one has to adjust the spectral zero level locally.  A number of \\mgtwo\\ systems show this effect in \\apm. By means of equivalent width modelling Petitjean et al. (2000) have obtained size estimates of $\\sim 1$~kpc for \\mgtwo\\ cloudlets. From separate near-infrared spectra for each LOS Kobayashi et al. (2002) obtain upper limit estimates of $\\sim 200$~pc for \\mgtwo\\ clouds in a damped \\lya\\ system.  We have used the ADR to select lines from intervening \\cfour\\ systems which may be due to coverage of one LOS towards \\apm.  The observations and data reduction are described in Section \\ref{sec:obs}, the fitting procedure in \\scr{sec:fit} and our partial coverage model in \\scr{sec:model}.  We have run a set of simulations to establish whether the effect can be genuine for each fitted line, and these are described in \\scr{simulation}. We have then used a maximum likelihood analysis to obtain estimates for the most probable sizes of the \\cfour\\ absorbers (Section \\ref{sec:size}). We have used the currently favoured cosmological model with matter density $\\Omega_M=0.3$, and vacuum energy density $\\Omega_{\\Lambda}=0.7$ (Perlmutter et al. 1999; Riess et al. 1998; Riess et al. 2001). To facilitate comparisons with previous work, where $q_0=0.5 \\ (\\Omega_M=1, \\Omega_{\\Lambda}=0)$ is used almost exclusively, we have also carried out the calculations for such a model. $H_0=72 \\ h_{72}$~\\kmps\\ Mpc$^{-1}$ (Freedman et al. 2001) throughout. We discuss our results in \\scr{sec:disc} and conclude in \\scr{sec:conc}.    \\section[]{Observations and data reduction}\\label{sec:obs}  The data used in this analysis were obtained through a programme of high resolution observations on the 10-m Keck I telescope in Hawaii in April and May 1998 using HIRES with the TeK 2048$\\times$2048 CCD at the Nasmyth focus. Details of the observations can be found in Ellison et al. (1999b). The data have been made publicly available (Ellison et al. 1999a) but we did not use the version available on the internet because there are a number of discontinuities due to inappropriate merging of the orders (also noticed by Petitjean et al. 2000). Instead, we used the original data files. However, in two of the exposures there was a spurious systematic shift in wavelength of $\\sim 0.2$~\\AA, due to incorrect dating and, hence, incorrect heliocentric velocity correction. After correcting for this problem, we used Tom Barlow's {\\bf MA}una {\\bf K}ea {\\bf E}chelle {\\bf E}xtraction programme (see http://www2.keck.hawaii.edu:3636/realpublic/inst/hires/m\\linebreak akeewww/index.html) to obtain a set of disjoint echelle orders, wavelength calibrated and rebinned onto a linear wavelength scale of 0.04~\\AA\\ per pixel.  We have flux calibrated each echelle order individually by means of the available standard stars and standard {\\small IRAF} routines to ensure that there were no discontinuities where they were joined.  The orders were joined by means of specially written interactive {\\small IRAF} scripts which ensured a smooth joining of adjacent and overlapping parts with weights proportional to the local signal-to-noise ratio (SNR).  A continuous spectrum has thus been formed for which the continuum does not change abruptly.  This continuum was fitted using cubic spline functions in spectral regions deemed free of absorption.  \\begin{figure} \\begin{center} {\\vskip-4mm} \\psfig{file=adr6452bw.ps, width=8cm} \\vspace{-4.0cm} \\end{center} {\\vskip-3mm} \\caption{Example of \\cfour\\ doublet showing an ADR.  The {\\it left} panel is for the 1548~\\AA\\ line and the {\\it right} panel for the 1550~\\AA\\ line. In each panel from top to bottom shown are data ({\\it solid} line) and fit ({\\it dotted} line), and residuals ({\\it solid} line) with $\\pm 1\\sigma$ level marked by two horizontal, {\\it dotted} lines. Flux units are arbitrary.} \\label{egzle} \\end{figure} ",
        "conclusions": "\\label{sec:conc} The spectra used in this work are the result of careful wavelength and flux calibration, and order merging. The observed ADRs are thus real, and we have shown that in many cases they can be simulated by means of a simple physical model. From LOS separations at different redshifts we obtained lower limit estimates of $\\sim 300 \\ h_{72}^{-1}$~pc for \\cfour\\ absorbers. From a statistical approach we obtained most probable sizes of a few kiloparsecs, but, because the sample is small, this estimate is rather uncertain.  We have briefly discussed the possible nature of the absorbers based on our size estimates. Clearly, more work needs to be done before we can decide among different possibilities. In this respect it would be interesting to carry out a cross-correlation study between Lyman break galaxies and our \\cfour\\ absorbers similar to the work by Adelberger et al. (2002) who found that the two classes may be related.  Ideally, one would like to have separate spectra for different LOS. However, given that lensed QSOs may have small image separations, it is difficult to find suitable candidates. The technique used here compensates for this problem as it allows probing of closer LOS. All that is required is the existence of two images with a known angular separation as well as an estimate for the redshift of the lens. As in the case studied here, the separation may be obtained in the course of different observations with different instruments (e.g. space based) or in different wavebands (e.g. radio) where the necessary resolution can be attained.  Obtaining a reliable value for $z_l$ is more difficult as it depends on a combination of modelling and detection of the lensing mass, but if there are several candidates, several models can be run.  In principle the method could be extended to more than two images. However, with several images of different brightnesses the uncertainties in the zero level adjustment parameter would probably allow different configurations. It may be possible to add a probability argument for these as well, but adding this degree of complexity for what would be a small sample is not likely to yield worthwhile results.  Finally, it is worth noting that ADRs can be useful in a different way as well.  In cases of spectra of QSOs which are not known to be lensed, the need for same, or similar, ADR corrections for several absorbers would be a strong indication that a QSO is lensed. The method illustrated in this paper would then have to be modified to include the angular separation as an unknown as well, with more speculative results."
    },
    "0212/astro-ph0212446_arXiv.txt": {
        "abstract": "{In an \\xmm raster observation of the bright local group spiral galaxy M33 we study the population of X-ray sources (X-ray binaries, supernova remnants, super-shells) down to a 0.5--10 keV luminosity of \\oergs{35} -- more than a factor of 10 deeper than earlier ROSAT observations. EPIC spectra and hardness ratios are used to distinguish between different source classes. We confirmed the 3.45 d orbital light curve of the X-ray binary M33 X7, detected a transient super-soft source in M33, and searched for short term variability of the brighter sources. We characterize the diffuse X-ray component that is correlated with the inner disk and spiral arms. We will compare the results with other nearby galaxies. ",
        "introduction": "The Local Group Sc galaxy M33 at a distance of 795 kpc (van den Bergh 1991) with relatively low inclination of 56$^{\\circ}$ (Zaritzky, Elston \\& Hill 1989) and its moderate Galactic foreground absorption (\\nh$ = 6\\,10^{21}$\\,cm$^{-2}$, Stark et al. 1992) is ideally suited to study the X-ray source population and diffuse emission in a nearby spiral. The Einstein X-ray Observatory detected diffuse emission from hot gas in M33 and 17 unresolved sources (Long et al. 1981; Markert \\& Rallis 1983; Trinchieri, Fabbiano \\& Peres 1988). First ROSAT HRI and PSPC observations revealed 57 sources and confirmed the detection of diffuse X-ray emission which may trace the spiral arms within 10' radius around the nucleus (Schulman \\& Bregman 1995; Long et al. 1996). Combining all archival ROSAT observations of the field, Haberl \\& Pietsch (2001; HP01 hereafter) found 184 X-ray sources within 50' radius around the nucleus, identified some of the sources by correlations with previous X-ray, optical and radio catalogues, and in addition classified sources according to their X-ray properties. They found candidates for super-soft X-ray sources (SSS), X-ray binaries (XRBs), supernova remnants (SNRs), foreground stars and active galactic nuclei (AGN) in the background.  Two M33 sources are known for their outstanding X-ray properties (Peres et al. 1989). The brightest source (X8, luminosity of about $10^{39}$\\,erg~s$^{-1}$) is the most luminous X-ray source in the Local Group of galaxies and coincides with the optical center of M33. Its time variability (Dubus et al. 1997), its point-like nature seen by ROSAT HRI (Pietsch \\& Haberl 2000) and Chandra (Dubus \\& Rutledge 2002) and its X-ray spectrum best described by an absorbed power law plus disk blackbody model (e.g. Ehle, Pietsch \\& Haberl 2001, E01 hereafter; La Parola et al. 2002) point at a black hole XRB. A possible periodicity of 106 days was not confirmed in later observations (see Parmar et al. 2001). The second source (X7) is an eclipsing XRB  with a binary period of 3.45 d and possible 0.31 s pulsations (Schulman et al. 1993, Dubus et al. 1997, 1999; D99 hereafter; Larson \\& Schulman 1997).  Here we present \\xmm raster observations which were carried out within the telescope scientist guaranteed time program to survey M33 homogeneously for point sources with a sensitivity of $10^{35}$\\,erg~s$^{-1}$ in the 0.5--10 keV band, a factor of ten deeper than previous surveys. First X-ray false color images combining different energy bands have been discussed in Pietsch (2002). ",
        "conclusions": "As demonstrated above, our setup of \\xmm raster observations of M33 yields very interesting results, fully confirming our expectations. We have detected more than 400 sources down to a luminosity limit which is below \\oergs{35} if the sources are located in M33. We showed the possibilities to classify them from their X-ray properties alone. There is unresolved emission detected at energies below 1 keV, most likely from hot gas in the inner region of the disk and possibly from the southern spiral arm and/or halo above, that needs to be followed up.  Several of the X-ray observations were heavily affected by high background and will be re-scheduled to homogenize the survey. Identifications of X-ray sources with surveys and follow-up observations at other wavelengths are in progress."
    },
    "0212/astro-ph0212500_arXiv.txt": {
        "abstract": "We present analytical expressions for the Fourier analog of the CMB three-point and four-point correlation functions, the spatial bispectrum and trispectrum, of the Ostriker and Vishniac effect in the linear and mildly nonlinear regime. Through this systematic study, we illustrate a technique to tackle the calculation of such statistics making use of the effects of its small-angle and vector-like properties through the Limber approximation. Finally we discuss its configuration dependence and detectability in the context of Gaussian theories for the currently favored flat $\\Lambda$CDM cosmology. ",
        "introduction": "\\label{sec:Introduction} In recent years, with the prospect of the increase in the sensitivity and angular resolution of the forthcoming Cosmic Microwave Background (CMB) satellite and interferometry experiments~\\cite{MAP/PLANCK,MINT,CBI,ACT}, efforts have been driven to the theoretical study of the secondary anisotropies contributions to the temperature fluctuations on arcminute scales and below. While the primordial anisotropies from recombination are thought to be well understood secondary anisotropies from reionization are not.  As is well known, the current favored inflationary model of structure formation predicts a nearly Gaussian probability distribution for the primordial density fluctuations. In this case, the CMB is completely described by its two-point correlation function or power spectrum in Fourier space. All higher-order correlations can be expressed in terms of it. Primordial nonlinearities and secondary effects introduce deviations from Gaussianity, producing a detectable signal in both the power spectrum and higher-order statistics. Recent work provides the theoretical background for the calculation of estimators of these higher-order statistics~\\cite{Hu01,Komatsu:Spergel00} and constrains possible non-Gaussian primordial contributions to the bispectrum and the trispectrum on degree and sub-degree angular scales using actual data ~\\cite{Kunz01,Santos01,Komatsu01,Troia03}. The interest is now in forecasting the expected signals on smaller scales due to secondary anisotropies, checking whether they are detectable and understanding how they can be separated from each other and from the primary anisotropies in light of the future data.  The Ostriker and Vishniac (OV) effect~\\cite{Ost:Vish86} was found to be the dominant linear secondary contribution to the CMB anisotropies below the Silk-damping scales at the arcminute level~\\cite{Hu94}. It is caused by Thomson scattering off of the CMB photons by moving electrons during the initial phase of reionization. It has the advantage of taking place during the linear regime of structure evolution and of being a small-angle effect enabling one to obtain analytical expressions for its higher-order correlation functions in the small-angle limit. Because of the highly predictive power of linear theory, any measurement of such statistics would be a sensitive probe of the reionization history of the universe, difficult to disentangle in measurements from nonlinear contributions. Several derivations for its power spectrum have been carried out~\\cite{Hu94,Dod:Jubas93,Vishniac87} but no one has fully addressed the calculation of its bispectrum or trispectrum.  It is then timely to obtain these expressions and to qualify and quantify them. We therefore extend and detail previous techniques used for the calculation of the OV power spectrum to the calculation of its higher-order statistics. As it will be shown, for the particular case of fields which are vector-like in nature, such as the OV effect, even moments will dominate over odd moments, making the trispectrum a much more sensitive statistics than the bispectrum.  Given the low redshifts of formation of structure, it is interesting to consider whether nonlinear effects can further enhance these statistics. So we will extend our study to the weakly nonlinear regime, allowing us to probe the most natural extension of the OV effect to nonlinear scales, the so-called kinetic Sunyaev-Zel'dovich (kSZ) effect. On small scales, both arise from the density modulation of the Doppler effect from large-scale bulk flows.  We review the relevant properties and parameters of the adiabatic cold dark matter (CDM) cosmology for structure formation in Sec.~\\ref{subsec:Cosm}. In Sec.~\\ref{subsec:OV} we review the theory of the OV effect and in Sec.~\\ref{subsec:Stat} we discuss the basic statistical properties of a general field through its n-point functions. In Sec.~\\ref{subsec:DomCont} we show how the homogeneous theory of turbulence combined with the Limber approximation enables one to infer the dominant contribution among n-point statistics of a vector field effect like the OV effect. In Sec.~\\ref{subsec:Signal-to-noise}, we introduce the standard formalism of the signal-to-noise calculation by means of the Fisher matrix. For general consistency, in Sec.~\\ref{subsec:PowSpectrum} we apply our method in detail to the calculation of the OV power spectrum as well as its nonlinear extension. In Sec.~\\ref{sec:Bispectrum} and in Sec.~\\ref{sec:Trispectrum} we present the steps of the calculation of the OV bispectrum and of the trispectrum respectively and its nonlinear counterparts. We also present the results. Finally in Sec.~\\ref{sec:Conclusions} we conclude. In the Appendix A we generalize the Limber approximation to the 3-point and 4-point correlation functions. This may be useful for other cosmological studies. ",
        "conclusions": "\\label{sec:Conclusions} Due to its strong predictive power, linear theory is a very sensitive probe of the early stages of the reionization history through the Ostriker and Vishniac effect. Analytical expressions for its correlation functions can be derived and their measurement would be of high value to our present knowledge of that still unclear epoch of the universe evolution. We have presented detailed calculations of the three Fourier statistics of interest of the OV effect, the power spectrum, the bispectrum and the trispectrum. For that purpose we have developed a new technique that allows one to obtain their dominant contributions under the Limber approximation framework. This method is applicable to the derivation of any statistics involving correlations among vector-like effects. It illustrates what was expected under statistical homogeneity and isotropy assumptions and the vector and small scale nature of the OV effect. We also evaluated numerically, as a function of scale and for a specific configuration (equilateral for the bispectrum and trapezoidal for the trispectrum), these statistics for a flat $\\Lambda$CDM cosmology and two reionization scenarios. The first one is based on our pre-WMAP knowledge ($z_r=8$) and the second one takes into account the high values for the electron optical depth measured by WMAP ($z_r=17$). We numerically obtained approximate scaling relations for the amplitude of the OV statistics on the reionization history considered and found that the dependence is stronger on the ionization fraction than on the redshift of reionization. We have also studied their detectability in view of future satellite experiments. The alternation of dominant/subdominant/dominant higher order correlation functions was numerically shown. While the bispectrum is undetectable even by a perfect multi-frequency experiment capable of subtracting the thermal SZ contributions, the trispectrum could be measured by Planck or by interferometer experiments targeting arcminute scales with high sensitivity and for a sufficiently long period of time. This provides a unique signal distinguishing the OV effect from other non vector-like secondary anisotropies and could be useful when trying to separate different physical mechanisms imprinting themselves on these measurable statistics. One should bear in mind that despite this useful characteristic signature, our results are quite optimistic, although encouraging, as they rely on various analytically helpful idealized assumptions, as described previously. Also, other contributions to the signal are to be expected at the arcminute level and thus further study of CMB small-scale secondary anisotropies and foregrounds contributions to the trispectrum is required and much justified. In order to obtain an upper limit on the possible kSZ contributions, we also extended our calculations to the mildly nonlinear regime. We found that, contrary to the bispectrum, there is a noticeable enhancement of the contributions of the trispectrum in a morphological dependent way and that this enhancement reflects itself on the calculations of the signal-to-noise. Hence nonlinearities are expected to enhance the even non-Gaussian signals produced by the OV effect and further complicate its disentanglement from inevitable model-dependent nonlinear effects arising from structure formation. Conversely, having a template of the OV effect can help in extracting the nonlinear contributions to reionization at those angular scales, providing another possible window on the complex physics associated with reionization."
    },
    "0212/astro-ph0212266_arXiv.txt": {
        "abstract": "We investigate the practical implementation of Taylor's (2002) 3-dimensional gravitational potential reconstruction method using weak gravitational lensing, together with the requisite reconstruction of the lensing potential. This methodology calculates the 3-D gravitational potential given a knowledge of shear estimates and redshifts for a set of galaxies. We analytically estimate the noise expected in the reconstructed gravitational field taking into account the uncertainties associated with a finite survey, photometric redshift uncertainty, redshift-space distortions, and multiple scattering events. In order to implement this approach for future data analysis, we simulate the lensing distortion fields due to various mass distributions. We create catalogues of galaxies sampling this distortion in three dimensions, with realistic spatial distribution and intrinsic ellipticity for both ground-based and space-based surveys. Using the resulting catalogues of galaxy position and shear, we demonstrate that it is possible to reconstruct the lensing and gravitational potentials with our method. For example, we demonstrate that a typical ground-based shear survey with redshift limit $z=1$ and photometric redshifts with error $\\Delta z=0.05$ is directly able to measure the 3-D gravitational potential for mass concentrations $\\ga10^{14} M_\\odot$ between $0.1\\la z\\la0.5$, and can statistically measure the potential at much lower mass limits. The intrinsic ellipticity of objects is found to be a serious source of noise for the gravitational potential, which can be overcome by Wiener filtering or examining the potential statistically over many fields. We examine the use of the 3-D lensing potential to measure mass and position of clusters in 3-D, and to detect clusters behind clusters. ",
        "introduction": "Gravitational lensing affords us a direct method to probe the distribution of matter in the universe, irrespective of its state or nature. This deflection of light by the gravitational potential of matter along its path can be observed as a local alteration of number counts of background galaxies (magnification), or a distortion of their shape (shear). It is the latter phenomenon which we will concern ourselves with here; we will further restrict ourselves to the case where this distortion is weak ($\\la 10$\\% change in the ellipticity of the object). Despite the weakness of the effect, and the intrinsic, nearly randomly orientated ellipticity of background galaxies, we can measure the weak shear by averaging the ellipticity or shear estimates of very many galaxies.  It has long been recognised that weak gravitational lensing is a valuable tool for examining the two dimensional projected matter distribution, and can consequently provide important information regarding large-scale structure (see e.g. Bartelmann \\& Schneider 2000, Bernardeau et al 1997, Jain \\& Seljak 1997, Kaiser 1998). In particular, the sensitivity of lensing to all the matter present, including the dominant dark matter, ensures that the weak lensing effect is an excellent probe for determining the quantity and distribution of matter.  Weak lensing studies for a wide range of galaxy clusters have been carried out, allowing precision measurements of the clusters' masses and mass distributions (see e.g. Tyson et al 1990, Kaiser \\& Squires 1993, Bonnet et al 1994, Squires et al 1996, Hoekstra et al 1998, Luppino \\& Kaiser 1997, Gray et al 2002). Moving to larger scales, the shear due to large-scale structure has been accurately measured by several groups (see e.g. van Waerbeke et al 2001, Hoekstra et al 2002, Bacon et al 2002, Refregier et al 2002, Brown et al 2002).  Redshift information has already been used in weak lensing studies, e.g. to determine the median redshifts of the lens and background populations; most analyses then project the lensing information into a 2-dimensional projected mass distribution. Wittman et al (2001, 2002) demonstrate the utility of using 3-D shear information by inferring the redshift of a cluster using shear and photometric redshift information for the galaxies in their sample. The importance of including redshift information to remove intrinsic galaxy alignments from shear studies has also been discussed (Heymans \\& Heavens 2002, King \\& Schneider 2002a,b). Lens tomography has been studied as a valuable means of introducing redshift information into shear power spectra (e.g. Seljak 1998, Hu 1999, 2002, Huterer 2002, King \\& Schneider 2002b), but only recently has the full reconstruction of the 3-D Dark Matter distribution from lensing been considered (Taylor 2002, Hu \\& Keaton 2002).  In this paper, we seek to discuss a practical implementation for reconstruction of the 3-D lensing and gravitational potentials from weak lensing measurements and redshift information (whether photometric or spectroscopic). We will use the reconstruction procedure of Taylor (2002), which allows us to calculate the entire 3-D gravitational potential if we have a knowledge of shear estimates and redshifts for a set of galaxies. This procedure is explained in Section 2.  In Section 3 we discuss sources of uncertainty for our reconstruction, examining analytically the effect of shot noise due to galaxy ellipticity, the effects induced by a finite survey, photometric redshift errors, redshift-space distortions, and multiple scatterings of light rays.  We aim to test the reconstruction method using simulations of realistic weak lensing data. In practice, for a particular volume of space containing a mass distribution, the data which would result from a real survey would be a set of galaxy ellipticities (in the weak lensing regime, these will be dominated by intrinsic ellipticity with a small gravitational lensing perturbation) and redshifts. The ellipticities will typically be defined by the galaxies' quadrupole moments (e.g. Kaiser et al 1995, Rhodes et al 2000) or estimated from a decomposition of galaxy shape into eigenfunctions (Refregier \\& Bacon 2002, Bernstein \\& Jarvis 2002). In order to provide a realistic catalogue of these data, we simulate mass distributions, calculate the expected lensing distortion in 3-D, and create a set of galaxies probing this distortion with a realistic redshift distribution and intrinsic ellipticity. We describe this procedure in detail in Section 4.1.  Given such a catalogue, we attempt to reconstruct the lensing potential and gravitational potential (Section 4.2). We use a generalised 3-D Kaiser-Squires (1993) inversion together with Taylor's (2002) formalism to obtain the 3-dimensional gravitational potential distribution.  In Section 5 we demonstrate the effectiveness of our implementation in reconstructing lensing and gravitational potentials. We examine the level of noise in our simulations in Section 6, including the Poisson noise associated with having only a finite number of galaxies probing the lensing distribution, and the noise from galaxies' intrinsic ellipticities. We go on to examine the uncertainties caused by photometric or spectroscopic redshift errors.  In Section 7 we study the utility of reconstructing only the 3-D lensing potential without continuing to the gravitational potential; this provides a useful method for detecting mass concentrations and measuring their mass and 3-D position. Finally, we summarise our results in Section 8. ",
        "conclusions": "In this paper, we have developed and tested a practical method for 3-D reconstruction of the gravitational potential via weak lensing measurements, together with the more easily obtained lensing potential. This methodology is based on the reconstruction equations of Kaiser \\& Squires (1993) and Taylor (2002), by which these local 3-D potentials can be calculated given knowledge of the lensing shear field in 3-D. This can be obtained by using shear estimators for galaxies with known redshifts.  We have presented analytical forms for the shot-noise uncertainty in the convergence, lensing potential, gravitational potential and density field, noting that these fields become progressively more noisy for realistic survey sizes. We have also calculated the effects of only having finite sky-coverage for a survey and estimated the variance measured on such a survey, and the additional uncertainty in the reconstruction due to structure beyond the survey boundary. In particular we have found that the contribution to the reconstruction uncertainty in the differential lensing potential, $\\Delta \\phi$, is dominated by large-scale structures.  We have also shown that further sources of error, including photometric redshift errors, redshift-space distortions, and multiple scatterings of light rays, will not be dominant in our reconstruction process.  In order to simulate the measurement of a 3-D gravitational field, we have calculated the expected lensing potential due to a given mass field upon a 3-D grid. From this we have calculated the shear expected upon a galaxy image for a galaxy positioned anywhere in the 3-D grid. We have produced a catalogue of galaxies positioned in accordance with a realistic redshift distribution and given each an appropriate shear plus a random intrinsic ellipticity; this final galaxy shear catalogue was the information given to our reconstruction software.  We have made reconstructions of the lensing and gravitational fields, by smoothing the noisy shear data, calculating the lensing potential according to Kaiser \\& Squires (1993) and using Taylor's equation (\\ref{inv}) to find the gravitational potential.  We have found that the method works well in reconstructing the full lensing and gravitational potentials in the absence of noise. However, the addition of Poisson sampling of the field at a finite set of galaxy positions, together with the noise due to the intrinsic ellipticity of objects, produces significant sources of error. We find that we obtain lensing potential maps with $\\nu \\simeq 6$ in [3', 3', 0.05] bins in angle and redshift for a cluster of mass $8\\times 10^{13}M_\\odot$. Unfortunately, corresponding gravitational potential maps have only $\\nu \\simeq 0.5$ in pixels of this size. However, applying Wiener filtering to this gravitational potential, we can obtain $\\nu \\simeq 3$ measurements of gravitational potential of clusters of mass $8\\times 10^{13}M_\\odot$.  This provides excellent prospects for obtaining cosmologically significant information directly from the measured gravitational potential field. For surveys with a redshift limit $z=1$, mass concentrations $\\ga 10^{14}M_\\odot$ can be directly mapped between $0.1 \\la z \\la 0.5$, while statistical mapping can be used to examine the gravitational potential fluctuations on smaller mass scales.  We have examined the effect of redshift uncertainties upon our simulations, and find that these contribute a much smaller error than the dominant intrinsic ellipticity noise term.  Finally we have emphasised that, even on scales where a gravitational potential measurement is uncertain, the 3-D measurement of the $\\phi$ field is valuable, and can give us useful information about the statistics and topology of the mass field. In particular, we can obtain accurate measurements of the mass and position of clusters along the line of sight, and significantly detect the presence of clusters behind clusters.  The methodology described here for engaging in 3-D gravitational mapping has numerous applications. We can obtain direct measurements of mass distributions in 3-D, which can act as an important cosmological probe via the mass function or cluster number counts. We can directly measure cross correlation functions between mass and light in 3 dimensions. Also, for possible 'dark' mass concentrations (e.g. Erben et al 2000), this 3-D mapping procedure allows us to measure both mass and radial position for such objects, which is quite impossible via conventional redshift methods."
    },
    "0212/astro-ph0212099_arXiv.txt": {
        "abstract": "Ultraluminous infrared galaxies (ULIRGs) are outstanding due to their huge luminosity output in the infrared, which is predominantly powered by super starbursts and/or hidden active galactic nuclei (AGN). NGC\\,6240 is one of the nearest ULIRGs and is considered a key representative of its class. Here, we report the first high-resolution imaging spectroscopy of NGC\\,6240 in X-rays. The observation, performed with the ACIS-S detector aboard the {\\sl Chandra} X-ray observatory, led to the discovery of two hard nuclei, coincident with the optical-IR nuclei of NGC\\,6240. The AGN character of {\\em both} nuclei is revealed by the detection of absorbed hard, luminous X-ray emission and two strong neutral Fe\\,K$\\alpha$ lines. In addition, extended X-ray emission components are present, changing their rich structure in dependence of energy. The close correlation of the extended emission with the optical H$\\alpha$ emission of NGC\\,6240, in combination with the softness of its spectrum, clearly indicates its relation to starburst-driven superwind activity. ",
        "introduction": "Ultraluminous infrared galaxies (ULIRGs) are characterized by their huge luminosity output in the infrared, which is predominantly powered by super-starbursts and/or hidden AGN (e.g., Genzel et al. 1998, Sanders 1999). Many distant {\\sl SCUBA} sources, massive and dusty galaxies, are believed to be ULIRG equivalents at high redshift (e.g., Barger et al. 1998, Lawrence 2001). Local ULIRGs, of which NGC 6240 is regarded a key representative, are therefore important laboratories to study the physics of superwinds driven by the nuclear starbursts, the processes of IGM enrichment, to search for  the presence of hidden AGN, and to study the physics of galaxy formation (many ULIRGs are interacting galaxies, believed to be on the verge of forming elliptical galaxies).  NGC\\,6240 is one of the nearest members of the class of ULIRGs, and has been intensely studied at virtually all wavelengths. The galaxy shows conspicuous loops and tails, and two optical nuclei separated by $\\sim$1.8$^{\\prime\\prime}$ (Fried \\& Schulz 1983) which are also detected in the IR and radio band. Recent observations suggest that NGC\\,6240 is in a relatively early merger state (Tacconi et al. 1999, Tecza et al. 2000, Gerssen et al. 2001). It is expected to finally form an elliptical galaxy. Ground-based optical spectroscopy of NGC\\,6240 shows LINER-like emission-line ratios.  Given the importance of the question whether an AGN is present, and whether it is the ultimate power-source of the ULIRG NGC\\,6240, efforts to search for such an AGN were undertaken at all wavelengths (see Sect. 1 of Komossa et al. 1998 for a review). Recent observations include the detection of flaring water-vapor maser emission from NGC\\,6240 (Nakai, Sato \\& Yamauchi, 2002), the presence of a high-excitation [OIV] infrared emission line (Genzel et al. 1998), and the compactness of the radio cores (Beswick et al. 2001). In X-rays, NGC\\,6240 exhibits exceptionally luminous extended emission (Schulz et al. 1998, Komossa et al. 1998, Kolaczyk \\& Dixon 2000) likely powered by superwinds, and an iron-line superposed on hard X-ray emission (Mitsuda et al. 1995, Schulz et al. 1998, Iwasawa \\& Comastri 1998) extending beyond 10 keV (Vignati et al. 1999, Ikebe et al. 2000), interpreted as arising from an absorbed active nucleus. It is the hard X-ray observations that give the best evidence for the presence of an AGN in NGC\\,6240.  It remained basically an open question, which of the two nuclei of NGC\\,6240 is the active one (or whether even both of them are); {\\sl HST} observations (Rafanelli et al. 1997) suggested that the southern nucleus harbors an AGN or a heavily obscured LINER.  The questions regarding the onset of starburst and AGN activity and their evolution in mergers are of fundamental importance for our understanding of AGN/black hole formation and evolution in general. Given the complex nature of the X-ray emission of NGC\\,6240 with contributions from many components suggested from previous X-ray observations (Schulz et al. 1998, Komossa et al. 1998, Iwasawa \\& Comastri 1998, Netzer et al. 1998, Vignati et al. 1999, Ikebe et al. 2000), spatially resolved spectroscopy is crucial in order for us to disentangle all contributing components, determine their nature, and derive their physical properties. In this {\\sl Letter}, we report the {\\sl Chandra} discovery that {\\em both} nuclei of NGC\\,6240 are active, and we show first results on the remarkably structure-rich extended X-ray emission.  Luminosities were calculated using a Hubble constant of $H_0 = 50$ km s$^{-1}$ Mpc$^{-1}$. At a distance of NGC\\,6240 of 144 Mpc, 1$^{\\prime\\prime}$ corresponds to 700 pc in the galaxy. ",
        "conclusions": "\\subsection{Starburst-related emission}  Could the widely extended X-ray emission of NGC\\,6240 still be powered by photoionization of the active nuclei ?  {\\sl Chandra} recently found some cases in which photoionization seems to dominate the ionization of circum-nuclear matter out to large distances ($\\sim$ 2\\,kpc in case of NGC\\,1068; Young, Wilson \\& Shopbell 2001). However, the correlation of X-ray and H$\\alpha$ emission in combination with the {\\em soft}, thermal-plasma-like spectrum of the extended emission of NGC\\,6240 strongly favors starburst-driven superwinds (Heckman, Armus \\& Miley 1987) as the origin for the bulk of the extended X-ray emission of NGC\\,6240. According to superwind models of Schulz et al. (1998), a mechanical input power of $L_{\\rm mech} = 3\\,10^{43}$ erg/s (derived from a SN-rate of 3/year) can, within 3\\,$10^7$ yrs, drive a single shell to an extent of $R \\sim$ 10\\,kpc within a medium of 0.1 cm$^{-3}$ particle density. This model reproduces the observed X-ray luminosity of the extended emission of $\\sim$10$^{42}$ erg/s. A distance of 10 kpc corresponds to the extent of the five-finger structure seen in H$\\alpha$, coincident with the brightest extended X-ray emission, including a correction for a disk inclined by up to 40$^{\\rm o}$ from edge-on.   The ionized iron line observed in addition to the neutral iron line may originate in `Cas-A-type' supernova remnants (e.g., Willingale et al. 2002), or in a near-nuclear `warm scatterer' that is ionized by the AGN (Komossa et al 1998).  The extended soft X-ray component and its link with non-X-ray morphological structures will be further discussed in a follow-up paper (Komossa et al., in prep.). In the following, we concentrate on the core region of NGC\\,6240.    \\subsection{Emission from the binary AGN}  Using {\\sl Chandra} ACIS-S, we found that {\\em both} nuclei of NGC\\,6240 are emitters of luminous hard X-ray emission, on which a strong neutral iron K$\\alpha$ line is superposed. These properties identify both nuclei of NGC\\,6240 as active. In particular, a strong, neutral Fe K$\\alpha$ line is not produced in a starburst-superwind, but originates from fluorescence in cold material illuminated by a hard continuum spectrum. For the first time, we can disentangle the contribution to the hard X-ray luminosity from the southern and northern nucleus. The observed (0.1-10)\\,keV, absorption-corrected, fluxes convert to {\\em observed} X-ray luminosities of the southern and northern nucleus, $L_{\\rm x, S}=1.9\\,10^{42}$ erg/s, $L_{\\rm x, N}=0.7\\,10^{42}$ erg/s, respectively.  The {\\em intrinsic} luminosities are significantly larger, because it has been repeatedly argued that the emission we are seeing in X-rays below 10 keV is scattered emission from the AGN (e.g., Mitsuda et al. 1995, Schulz et al. 1998, Komossa et al. 1998). The intrinsic AGN emission of NGC\\,6240 only shows up above $\\sim$9-10 keV (Vignati et al. 1999, Ikebe et al. 2000). The strength of the neutral Fe K$\\alpha$ line in both nuclei suggests a scattering geometry for {\\em both} AGN. This is also indicated by the flatness of the hard powerlaw spectra (Sect. 4.2).  We find marginal evidence that part of the hard emission is extended by few arcsesonds. 3$^{\\prime\\prime}$ correspond to $\\sim$2 kpc in NGC\\,6240. It is interesting to note that this scale is similar to the extent of the hard X-ray emission and iron line of  NGC\\,1068 detected with {\\sl Chandra} (Young, Wilson \\& Shopbell 2001). Dense cold gas that could provide the scattering agent is abundant in the center region of NGC\\,6240. Deeper {\\sl Chandra} observations of NGC\\,6240, still employing the superb imaging spectroscopy capabilities of the ACIS-S detector, will be indispensable for studying the complex central region of NGC\\,6240 in greater detail; for instance to obtain a high-quality image in the Fe\\,K$\\alpha$ line.  The presence of supermassive binary black holes, as in NGC\\,6240, are of importance for our understanding of the formation and evolution of AGN and the formation of elliptical galaxies via mergers.  Ultimately, the binary AGN of NGC\\,6240 will coalesce to form one nucleus. The final merging of the supermassive black holes is expected to produce a strong gravitational wave signal. In fact, such events are expected to generate the clearest signals detectable with the gravitational wave detector LISA that will be placed in Earth orbit in the near future (e.g., Danzmann 1996)."
    },
    "0212/astro-ph0212320_arXiv.txt": {
        "abstract": "To understand galaxies and their evolution, it is necessary to describe how the different scales interact: how the microscopic physics, such as star formation, or the large scale physics, such as galaxy interactions may modify the galaxy global shapes.  The purpose of this review is to point out some general or recent topics related to such scale interactions, both observational and theoretical, which are relevant in the present understanding of galaxies. ",
        "introduction": "The purpose of this talk is to review the general mechanisms acting between the scales in gravitating systems such as galaxies, and that lead to evolution.  The Universe is certainly hierarchically organized, and at different scales different phenomena interact with the adjacent scales.  For the sake of simplicity it has been customary for a long time to consider one scale at a time and to neglect scale interactions.  Scientists have been inclined to discard scale interactions already because distinct disciplines are specialized according to the astrophysical object sizes.  As a consequence some attempts to introduce scale interactions such as galaxy mergers met much resistance.  But nowadays it becomes clearer that as stars are built upon nuclear and atomic physics, galaxies are also built upon the ISM physics and star formation processes.  Conversely galaxy clusters and large scale structures do matter for understanding galaxy evolution, which in turn cannot be ignored when trying to explain star formation, while stellar evolution do certainly plays a role in the chemical composition of matter.  Despite a disparity of phenomena, we see common phenomena through the scales because some of the involved physics is scale-free, namely gravitational dynamics.  Statistical physics provides an example how adjacent scales may be sometimes absorbed by a proper theoretical frame.  The small scale physics is absorbed by statistical concepts, such as the ``molecular chaos'', while the large scale physics is absorbed by proper boundary conditions, such as the extensivity of the system.  Thus one can obtain a useful description of gases without including the detailed knowledge of molecules, which are actually quite complex objects. Without any particle computer simulations statistical physics can describe the fate of a gas enclosed in a box, the shape of which is unimportant.  However, if the box size increases too much we know well that a sufficiently large volume of gas becomes self-gravitating, and then the usual tools of statistical mechanics do not work well with gravitating systems.  We briefly review on the reasons in Sect.~2.  In Sect.~3 we discuss the small scale interactions in galaxies, i.e., the ISM and star formation physics for which very interesting developments have occurred in the recent years.  In Sect.~4 some rarely discussed issues about the large scale interactions of galaxies are presented. Finally, we conclude in Sect.~5. ",
        "conclusions": "Scale interactions are important in astrophysics because the negative specific capacity states are ubiquitous.  These states are unstable and develop phase-space correlations which trigger the interactions of scales.  This poses several deep problems for the current way to perform N-body simulations, but this also indicates that galaxies should not be seen as isolated system insensitive to their scale boundary conditions.  At small scale stellar mass loss is secularly important, over several Gyr galaxies must be seen as dissipative structures.  Also the stellar energy input is important enough to compete with dynamics and is sufficient to determine the global spiral parameters.  At large scales, galaxy interactions lead to a a walk through the mass, energy, and angular momentum space, since at each interaction a substantial amount of matter may be recycled through the IGM.  Also the angular momentum constraint linked to dark matter rich gas infall indicates that flaring disk-like rotationally supported dark matter distributions should be considered too, not exclusively pressure supported spheroids."
    },
    "0212/astro-ph0212116_arXiv.txt": {
        "abstract": "{ We discuss the first part of a survey for RR Lyrae variables in the galactic halo that is being made with the 1m Schmidt telescope at the Venezuelan National Observatory.  So far the survey has discovered 497 variables in 380 deg$^2$, lying from 4 to 60 kpc from the Sun.  It has detected three statistically significant clumps of variables, which shows that outer halo does not have smooth density contours.  One clump is located at 50 kpc from the galactic center, and it is probably tidal debris from the Sagittarius dwarf spheroidal galaxy.  A second structure, at 17 kpc from the galactic center, appears to be due to the tidal disruption of the globular cluster Pal 5.  The third group, at 19 kpc, is not related with any known globular cluster or dwarf galaxy.  The little that is known about its properties is consistent with it being debris from a disrupted galaxy. ",
        "introduction": "Wide field imagers are essential tools for modern studies of the large-scale structure of the halo of our Galaxy since it is necessary to observe hundreds of square degrees of the sky. A second key ingredient is to use a tracer of the halo population that is easily separated from the numerous foreground stars belonging to the thin and thick disk populations. Among several other good possibilities, variable stars of the RR Lyrae type are ideal tracers. They are low-mass evolved stars and therefore, they trace the oldest stellar population (age $> 9$ Gyrs). They are well known standard candles, which make them ideal for studying the three-dimensional structure of the halo. And finally, they are very easy to recognize because of their large-amplitude variability ($\\sim 1$ mag) and relatively short periods.  In recent years our understanding of the structure of the halo has advanced significantly. It is now clear that at least the outer regions of the halo, do not have a smooth distribution of stars. Instead, the distribution seems to be quite clumpy \\citep{new02,viv01,yan00,ive00}. The most likely interpretation is that these sub-structures are relics of small satellite galaxies that have been accreted and destroyed by the tidal forces of the Milky Way. The stars (and globular clusters) of those disrupted satellite galaxies have formed a part, if not all of the halo of our Galaxy. This scenario is also consistent with theoretical models of hierarchical formation of structures based in cold dark matter cosmology \\citep[eg.,][]{moo99,bul01}.  In order to quantify how important was the accretion mechanism in the formation of the halo, we need to find and study the relics of these disrupted galaxies. Moreover, their study may reveal a relationship with the population of dwarf spheroidal (dSph) galaxies that still orbit the Galaxy. The ultimate goals are to determine the accretion history of the Galaxy and the fates of its retinue of past and present satellite galaxies.  We present here the first part of a large scale survey of RR Lyrae variable stars (RRs) aimed at studying the spatial distribution of stars in the halo, especially at large distances from the galactic center. This distant region of the Galaxy has been poorly studied in the past mainly because of the lack of wide field CCD cameras able to cover large regions of the sky, down to magnitudes faint enough to reach the far halo. ",
        "conclusions": "The streams and sub-structures found in the distribution of RRs in the outer halo of the Milky Way resemble the signatures of merger events. The outer halo is significantly clumpy and this is consistent with the idea that accretion of satellites played a major role in the formation of the outer halo.  Satellite dSph galaxies are the main suspect for being the building blocks of the halo. If so, the stellar population of the halo should resemble a composite of several dSph galaxies. The pulsational properties of the halo and these galaxies are indeed similar in that both have wide period-amplitude distributions and mean periods intermediate between Oo groups I and II. Although this may not constitute a strong proof that the halo RR Lyrae population came indeed from disrupted dSph galaxies, it is at least consistent with this hypothesis.  Our data have confirmed previous studies that showed that the inner halo is flattened towards the galactic plane.  This may be evidence that the inner halo formed differently, perhaps as a dissipative collapse of a protogalactic cloud.  The search for streams in the halo by QUEST and other surveys like the SDSS still cover a small fraction of the sky. Several streams have been found but they are still few to allow a description of the accretion history of the Milky Way. We are continuing the RR Lyrae survey with the QUEST camera in a second declination strip. An important by-product of this survey will be a very large database of variable objects which will be made available to the astronomical community in the near future."
    },
    "0212/nucl-th0212093_arXiv.txt": {
        "abstract": "The bottom part of the neutron star crust is investigated using the Skyrme-Hartree-Fock approach with the Coulomb interaction treated beyond the Wigner-Seitz approximation. A variety of nuclear phases is found to coexist in this region. Their stability and relative energies are governed by the Coulomb, surface and shell energies. We have also found that a substantial contribution is coming from the spin-orbit interaction. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/astro-ph0212428_arXiv.txt": {
        "abstract": "The CZT detector on the \\infocus\\ hard X-ray telescope is a pixellated solid-state device capable of imaging spectroscopy by measuring the position and energy of each incoming photon.  The detector sits at the focal point of an 8\\,m focal length multilayered grazing incidence X-ray mirror which has significant effective area between 20--40 keV. The detector has an energy resolution of 4.0\\,keV at 32\\,keV, and the \\infocus\\ telescope has an angular resolution of 2.2 arcminute and a field of view of about 10 arcminutes.  \\infocus\\ flew on a balloon mission in July 2001 and observed Cygnus X-1.  We present results from laboratory testing of the detector to measure the uniformity of response across the detector, to determine the spectral resolution, and to perform a simple noise decomposition.  We also present a hard X-ray spectrum and image of Cygnus X-1, and measurements of the hard X-ray CZT background obtained with the SWIN detector on \\infocus. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/astro-ph0212102_arXiv.txt": {
        "abstract": "Low Surface Brightness (LSB) galaxies are dominated by dark matter, and their rotation curves thus reflect their dark matter distribution. Recent high-resolution rotation curves suggest that their dark matter mass-density distributions are dominated by a constant-density core. This seems inconsistent with the predictions of Cold Dark Matter (CDM) models which produce halos with compact density cusps and steep mass-density profiles.  However, the observationally determined mass profiles may be affected by non-circular motions, asymmetries and offsets between optical and dynamical centres, all of which tend to lower the observed slopes.  Here we determine the impact of each of these effects on a variety of halo models, and compare the results with observed mass-density profiles.  Our simulations suggest that no single systematic effect can reconcile the data with the cuspy CDM halos. The data are best described by a model with a soft core with an inner power-law mass-density slope $\\alpha = -0.2 \\pm 0.2$. However, no single universal halo profile provides a completely adequate description of the data. ",
        "introduction": "Low Surface Brightness (LSB) galaxies are dominated by dark matter. Their stellar populations are dynamically unimportant and the rotation curves are thus direct tracers of their dark matter distributions \\citep{edb_rot}. This is important, as in brighter galaxies it is difficult to disentangle the dynamical contributions of the dark and visible matter, and unambiguously determine the distribution of the dark matter. LSB galaxies give us a unique opportunity to compare observations with cosmological dark matter simulations.  These numerical models, based on the Cold Dark Matter (CDM) paradigm, make very specific predictions about the distribution of dark matter in galaxies \\citep{dubinskicdm}. In CDM halos the mass density distribution in the inner parts is characterized by a steep density ``cusp'', usually described by a power-law $\\rho(r)\\sim r^{\\alpha}$ with slopes varying from $\\alpha = -1$ \\citep{NFW96,NFW97} to $\\alpha = -1.5$ \\citep[e.g.][]{moore98,moore99,bul99}. This cusp manifests itself in a steeply rising rotation curve with a specific shape. Observations of dwarf and LSB galaxies \\citep{moore94, optcur_data, optcur_pap2, blais01, lsbopt_bosma, marchesini02} show that real rotation curves rise less steeply than predicted, and do not have the required CDM shape. Similar conclusions have been reached by \\citet{salucci01} and \\citet{salbor} for high surface brightness disk galaxies.  At small radii, the empirically determined mass distribution can be well described by a central kpc-sized constant-density core (i.e.\\ $\\alpha \\sim 0$ in the inner parts) (\\citealt{optcur_letter,lsbopt_bosma} [hereafter dBMBR and dBB02 respectively]).  dBMBR and dBB02 base their conclusions on observations of rotation curves of over 60 LSB galaxies. They determine the mass density profiles that give rise to the observed rotation curves and fit the inner parts with a single power-law. The resulting distribution of inner mass-density slopes shows a strong peak at $\\alpha = -0.2 \\pm 0.2$ (see Fig.~2 in dBMBR), but with a large spread, varying from steep and cuspy $\\alpha=-2$ profiles to flat and even moderately positive slopes.  This was interpreted as a result of resolution effects: dBMBR find that the best resolved curves show the flattest profiles (see Fig.~3 in dBMBR and Fig.~14 in dBB02), contrary to what is expected for CDM.  At lower resolutions the CDM and the ISO models both predict the same slopes, leading to an ambiguity as to which model fits lower resolution data best \\citep{vdB&S}.  However, resolution is not the only effect that can influence the measured slopes. Other effects, which dBMBR and dBB02 did not thoroughly investigate, include non-circular motions, offsets between dynamical and optical centres, lopsidedness, asymmetries, and slit-width.  These effects all have the potential to decrease the measured inner mass-density slopes. It is thus important to investigate if the observed shallow slopes truly reflect the underlying mass distribution, or whether they are determined by systematic effects.  For example, small offsets between the dynamical centre and the optical centre are sometimes observed in late-type barred galaxies \\citep{DJ98,deVauc72}. If this were a systematic property of all LSB or dwarf galaxies then observations aligned on the optical centre would not necessarily probe the dynamical centre.  Even if the galaxy in question had a steep CDM density cusp at the dynamical centre, the derived slope would be shallower, giving the impression of a core-dominated halo. Hence these systematic effects should be understood, in order to judge whether LSB galaxies have cuspy or core-dominated halos.  Furthermore, once the impact (if any) of these systematic effects is understood, one could hope to determine how much they contribute to the scatter in the observed slopes, and whether the range in observed slopes can be attributed to ``cosmic scatter'' or whether a single halo profile for all galaxies suffices.  In this paper we investigate the importance of these systematic effects by deriving a large ensemble of rotation curves of both cuspy CDM and core-dominated (simulated) halos, using observed curves as a constraint.  We apply the relevant systematic effects to these simulated curves, and then determine the inner mass-density slopes. Our main conclusion is that none of the systematic effects investigated can convincingly wipe out the signature of a CDM halo. In other words, if LSB galaxies had CDM halos, they would be unambiguously detectable, even in the presence of systematic effects. The most likely interpretation is that the observed core-dominated mass-distribution reflects the true dark matter distribution in LSB galaxies.  The organization of this paper is as follows.  Section~2 discusses some aspects of the observational comparison sample. Section~3 describes the simulations.  Section~4 contains the results of simulations that include various systematic effects. In Sect.~5 we describe some tests performed with real data, while in Sect.~6 we attempt to use the simulations to improve on the halo model used. In Sect.~7 we present our conclusions. ",
        "conclusions": "We have investigated the effect of systematic effects such as non-circular motions, asymmetries and lopsidedness on the values of the inner slopes of various halo models.  We find that no realistic combination of a single systematic effect and NFW models can explain the observed distribution of slopes.  Core-dominated models, on the other hand, are consistent with the data in the absence of large systematic effects.  The best match to the data is obtained by using halos with an inner mass-density slope $\\alpha=-0.2$ and an intrinsic Gaussian scatter in the slope of $\\sigma_{\\alpha} \\sim 0.2$. As our models assume minimum disk, these slopes should be interpreted as those of the combined mass-profiles of stars, gas and dark matter in a galaxy.  It is perhaps possible to construct a combination of multiple systematic effects to explain the data in the CDM framework.  However, such a combination of effects would have to completely wipe out the strong $\\alpha = -1$ peak in the $\\alpha$ histograms. The resulting models would be completely dominated by non-cosmological effects, and not contain any link between properties of dark matter halos and the visible galaxies that inhabit them.  Yet, they would have to explain and obey the Tully-Fisher relation, for example, and thus require a tremendous amount of fine-tuning.  The implied observational signature of CDM halos is strong and if present should be easily seen.  We thus conclude that the trends observed in the $r_{\\rm in} - \\alpha$ plot are mostly resolution effects, in combination with intrinsic scatter in the inner slopes.  The cusp problem is certainly genuine; it can not plausibly be attributed to systematic errors in the data. LSB galaxies have shallow mass-density slopes implying that (i) the current CDM model, or more precisely the current generation of N-body models based on the CDM prescription, do not correctly describe structure at the scale of galaxies, or (ii) non-cosmological effects destroy the cusp.  The second option reduces the elegance and predictive power of the CDM model."
    },
    "0212/astro-ph0212334_arXiv.txt": {
        "abstract": "{We  present UBVRIZJsHKs  broad band  photometry of  the host galaxy  of the  dark gamma-ray  burst  (GRB) of  February 10,  2000. These observations represent the most exhaustive photometry given to date of any GRB host galaxy.  A grid of spectral templates have been fitted to the  Spectral Energy Distribution (SED) of  the host.  The derived  photometric redshift is  $z=0.842^{+0.054}_{-0.042}$, which is   in  excellent   agreement  with   the   spectroscopic  redshift ($z=0.8463\\pm  0.0002$)  proposed by  Piro  et al.\\  (\\cite{Piro02}) based on  a single emission  line.  Furthermore, we  have determined the  photometric  redshift  of  all  the  galaxies  in  an  area  of $6^{\\prime} \\times  6^{\\prime}$ around the host galaxy,  in order to check for their overdensity in the environment of the host.  We find that the GRB~000210 host galaxy is a subluminous galaxy ($L \\sim 0.5 \\pm  0.2  L^{\\star}$),  with   no  companions  above  our  detection threshold  of $0.18  \\pm 0.06  L^{\\star}$.  Based  on  the restframe ultraviolet flux  a star formation  rate of $2.1 \\pm  0.2 M_{\\odot}$ yr$^{-1}$ is estimated.   The best fit to the SED  is obtained for a starburst template with an  age of $0.181^{+0.037}_{-0.026}$ Gyr and a  very  low  extinction  ($A_{\\rm  V} \\sim  0$).   We  discuss  the implications of the inferred low value of $A_{\\rm V}$ and the age of the  dominant  stellar  population  for  the non  detection  of  the GRB~000210  optical  afterglow. ",
        "introduction": "The origin of cosmological Gamma-Ray  Bursts (GRBs) remains one of the great    mysteries   of    modern   astronomy    (van    Paradijs   et al. \\cite{vanP00}).  Over the past half decade many advances have been made in  understanding the nature  of the bursts and  their afterglows throughout the electromagnetic spectrum.   There are at present mainly two sets  of models for  GRBs.  One set  of models predicts  that GRBs occur  when two  collapsed objects  (such  as black  holes or  neutron stars)  merge  (Eichler  et  al.  \\cite{Eich89};  Mochkovitch  et  al. \\cite{Moch93}).  The time-scale for binary compact objects to merge is large  ($\\gtrapprox$ 1  Gyr), so  GRBs  can occur  after massive  star formation  has ended  in  a galaxy.   The  other major  set of  models predicts  that GRBs  are associated  with the  death of  massive stars (supernovae   or  hypernovae)   (Woosley   \\cite{Woos93};  Paczy\\'nski \\cite{Pacz98}; MacFadyen \\& Woosley \\cite{MacF99}).  In this case GRBs will  coincide with  the  epoch of  star  formation in  the host.   By determining the Spectral Energy  Distribution (SED) and star formation rate (SFR) of a sample of GRB host galaxies we can distinguish between these two  families of GRB  progenitor models (see also  Belczynski et al. \\cite{Belc02}).  Substantial insight has already been gained about the galaxies that the bursts occur in.  Radio, optical and/or infrared afterglows have been  observed for $\\sim$40 GRBs, and  the majority of these coincide with starforming galaxies.  As GRB host galaxies tend to be faint (R$ > 23$) spectroscopic studies of  the SED  are  only reachable  with  8--10 m  class telescopes.   A cheaper  and  elegant  alternative   to  spectroscopy  is  to  extract information  on   the  properties  of  the  host   galaxies  based  on multicolour  broad band imaging.   By determining  the colours  of GRB host galaxies  we can derive or  constrain the age  of the predominant stellar population  as well as the  extinction. As part  of the global fit, the photometric  redshift of the host galaxies  can be derived if the redshift  is not known in advance  from spectroscopic observations of the afterglow and/or the host galaxy.  Additional advantages of the multicolour photometric  studies compared to  spectroscopic techniques are  their   simplicity  and  their   multi-object  feasibility.   The photometric technique  allows the determination of the  colours of all objects  in the  field  down to  the  imaging flux  limit, thereby  in principle  permitting the study  of the  host galaxy  environment. The precision of  the photometric redshift estimate (which  depends on the photometric accuracy,  the spectral coverage and the  number of bands) is  evidently not  as accurate  as the  spectroscopic one,  but  it is sufficient for a first order study of host galaxy environments.  So  far it has  been possible  to detect  optical afterglows  for only about 30\\% of localised GRBs (Fynbo et al.  \\cite {Fynb01}; Lazzati et al.  \\cite{Lazz02}).  It is important  to understand the nature of the remaining (rather ill-termed) so-called dark  GRBs if we wish to get a complete understanding on GRB  selected galaxies and thereby constrain the  GRB  progenitors as  well  as  the  distribution of  cosmic  star formation    over    different    modes   (e.g.     Ramirez-Ruiz    et al. \\cite{Rami02};  Venemans \\&  Blain \\cite{Vene01}).  GRB~000210 is currently one  of only few  systems that allow  a detailed study  of a galaxy  hosting   a  dark  GRB.   The  burst   exhibited  the  highest $\\gamma$-ray peak flux  among the 54 GRBs localized  during the entire BeppoSAX operation,  from 1996 to  2002 (Piro et  al.  \\cite{Piro02}). However, no  optical afterglow  (OA) was detected  in spite of  a deep search (R$> 23.5$) carried out  $\\sim16$ hrs after the gamma-ray event (Gorosabel et al.  \\cite{Goro00a}).  X-ray observations performed with the Chandra X-ray  telescope 21 hrs after the  GRB localised the X-ray afterglow of  the burst to  an accuracy of $2^{\\prime  \\prime}$, later improved by Piro et al.   (\\cite{Piro02}) to a $0\\farcs6$ radius error circle.   The  optical search  revealed  an  extended constant  source coincident with the X-ray afterglow which was proposed as the GRB host galaxy (Gorosabel  et al.  \\cite{Goro00b}).  In addition,  Piro et al. (\\cite{Piro02}) have  reported the detection  of a radio  transient at 8.5 GHz spatially  coincident with the X-ray afterglow.   Based on the detection of a single host galaxy spectral line, interpreted to be due to [\\ion{O}{II}], Piro et  al.  (\\cite{Piro02}) proposed a redshift of $z=0.8463 \\pm  0.0002$.  Recently  Berger et al.   (\\cite{Berg02}) and Barnard  et  al.   (\\cite{Barn02})  have reported  $\\sim$2.5  $\\sigma$ detections  of sub-mm  emission  towards the  position  of GRB~000210 interpreted as  emission from the  host galaxy and hence  suggesting a SFR of several hundred $M_{\\odot}$ yr$^{-1}$.  \\begin{table*} \\begin{center} \\caption{Chronologically ordered optical  and NIR observations carried out for the GRB~000210 host galaxy.} \\begin{tabular}{lcccccr} Telescope     &  Filter & Date UT &T$_{\\rm exp}$&  Seeing  & Limiting magnitude \\\\ (+Instrument) &         &         &  (s)       &          &    ($4\\sigma$)     \\\\ \\hline 8.2VLT (+FORS1)  &  R  &  25.237--25.240/10/00&           300 & 0.70& 25.4$^{\\star\\star}$&\\\\ 3.58NTT (+SOFI)   &  H  &  02.251--02.410/09/01& 182$\\times$60 & 0.90& 22.8\\\\ 3.6ESO (+EFOSC2) &  V  &  13.219--13.253/09/01&   4$\\times$600& 1.75& 25.4\\\\ 3.6ESO (+EFOSC2) &  I  &  13.256--13.278/09/01&   3$\\times$600& 1.45& 23.1\\\\ 3.6ESO (+EFOSC2) &  B  &  13.280--13.302/09/01&   3$\\times$600& 1.70& 25.6\\\\ 3.6ESO (+EFOSC2) &  U  &  13.304--13.348/09/01&   6$\\times$600& 1.55& 24.7\\\\ 8.2VLT (+ISAAC)  &  Ks &  21.159--21.193/09/01&  30$\\times$60 & 0.45& 22.2\\\\ 8.2VLT (+ISAAC)  &  Js &  21.194--21.218/09/01&  15$\\times$120& 0.60& 24.1$^{\\dag}$ \\\\ 8.2VLT (+ISAAC)  &  Js &  23.193--23.218/09/01&  15$\\times$120& 0.75& 24.1$^{\\dag}$ \\\\ 1.54D  (+DFOSC)  &  Z  &  19.090--19.254/12/01&  14$\\times$600& 1.10& 22.9$^{\\star}$ \\\\ 1.54D  (+DFOSC)  &  Z  &  20.042--20.394/12/01&  21$\\times$600& 1.15& 22.9$^{\\star}$ \\\\ \\hline \\multicolumn{7}{l}{$\\star\\star$     Published in Piro et al. (\\cite{Piro02}).}\\\\ \\multicolumn{7}{l}{$\\dag$     The images were coadded resulting in just a single Js-band magnitude.}\\\\ \\multicolumn{7}{l}{$\\star$    The images were coadded resulting in just a single Z-band magnitude.}\\\\ \\hline \\label{table1} \\end{tabular} \\end{center} \\end{table*}  In this paper  we present the most intensive  multi-colour host galaxy imaging performed to date.  The  host galaxies SED studies to date had a limited number of bands (Sokolov et al.  \\cite{Soko01}; Chary et al. \\cite{Char02})    and   no   photometric    redshift   determinations. Throughout, the  assumed cosmology  will be $\\Omega_{\\Lambda}  = 0.7$, $\\Omega_{M} =  0.3$ and  $H_0= 65$ km  s$^{-1}$ Mpc$^{-1}$  (except in Sect.  \\ref{subluminous} where the  host galaxy luminosity is rescaled to  the  cosmology used  by  Lilly  et  al.  \\cite{Lill95}).   At  the proposed spectroscopic  redshift ($z=0.8463$),  the look back  time is 7.59 Gyr  (52.4\\% of the present  age) and the  luminosity distance is $1.79 \\times 10^{28}$ cm.  The  physical transverse size of one arcsec at $z=0.8463$ corresponds to 8.24 kpc. ",
        "conclusions": "We have  presented an intensive UBVRIZJsHKs broad  band photometry of the  GRB~000210 host  galaxy which  has allowed  us to  determine its photometric   redshift.    The   derived  photometric   redshift   is $z=0.842^{+0.054}_{-0.042}$,   in   excellent   agreement  with   the spectroscopic  redshift ($z=0.8463\\pm  0.0002$) proposed  by  Piro et al.\\ (\\cite{Piro02}).  The inferred redshift is basically independent of  the extinction law  and IMF  assumed, although  (at least  in the particular case of GRB~000210) the Scalo (\\cite{Scal86}) IMF provides slightly  worse  results than  Miller  \\&  Scalo (\\cite{Mill79})  and Salpeter  (\\cite{Salp55})  IMFs.   The  SED  of the  host  galaxy  is consistent with a  starburst template with an age  of $\\sim$0.181 Gyr and  a  very low  extinction  ($A_{\\rm V}  \\sim  0$).   Based on  the restframe UV  flux a star formation  rate of $2.1  \\pm 0.2 M_{\\odot}$ yr$^{-1}$ is estimated.  The absolute restframe B-band magnitude  of the host ($M_B = -20.16$) is consistent with  the distribution of the $M_B$  host galaxy values measured to date (see Djorgovski  et al.  \\cite{Djor01}, Fig. 2).  We determine a value of $L  = 0.5\\pm0.2 L^{\\star}$ for the luminosity of the  host, in  agreement  with the  value  estimated by  Piro et  al. (\\cite{Piro02}).  We have tried to the explore the role played by galactic interactions triggering the  GRB phenomena.  Many  host galaxies observed  to date appear as part of complex and interacting systems (GRB~980613, Hjorth et   al.   \\cite{Hjor02};   GRB~001007,  Castro   Cer\\'on  et   al. \\cite{Cast02}).  According to our study the GRB~000210 host galaxy is a subluminous  galaxy with no interacting companions  above $0.18 \\pm 0.06 L^{\\star}$.  The low value  of the extinction obtained in the  SED fit ($A_{\\rm V} \\sim 0$)  makes  difficult   to  explain  the  optical  darkness  of GRB~000210 in  terms of  the global host  galaxy dust  extinction. If dust  extinction is  the  reason  of the  lack  of optical  afterglow emission,  then the  circumburst region  has to  be very  compact and localised around  the progenitor.   This hypothesis would  agree with observations   carried  out   for  the   optically-faint  GRB~990705. Andersen  et al.   (\\cite{Ande02}) have  localised the  optically dim GRB~990705 (but  NIR bright, see  Masetti et al. \\cite{Mase00})  in a face-on spiral galaxy.   Thus given the thin disk  of a spiral galaxy ($\\sim 0.3$  kpc), the  optical extinction of  GRB~990705 can  not be attributed to the global ISM in its host.  This clumpy and fragmented ISM  would also  explain  the apparent  discrepancy  between our  SFR estimate  (derived from  the galaxy  UV  flux) and  the one  recently reported   based  on   the  sub-millimeter   range  (Berger   et  al. \\cite{Berg02}; Barnard et al.  \\cite{Barn02}).  Several progenitor models have been discussed in order to explain the inferred stellar  population age and the low  host galaxy extinction. Both  the  collapsar  and  the  binary  merging  models  show  severe limitations to explain the visible  stellar age and the line of sight \\ion{H}{I} column density (derived from the afterglow X-ray spectrum) respectively.  A solution to this problem would be the existence of a younger population of stars (several Myr of age) hidden by the clumpy ISM.   Such population  (which would  include the  progenitor massive star)  would   not  have   any  impact  in   the  host   galaxy  SED. Morphological information  derived by  HST could verify  the proposed ISM  clumpy  scenario  present  in   the  host  galaxy  of  the  dark GRB~000210."
    },
    "0212/astro-ph0212514_arXiv.txt": {
        "abstract": "{It is widely believed that gamma-ray bursts are produced by a jet-like outflows directed towards the observer, and the jet opening angle($\\theta_j$) is often inferred from the time at which there is a break in the afterglow light curves. Here we calculate the GRB afterglow light curves from a relativistic jet as seen by observers at a wide range of viewing angles ($\\theta_v$) from the jet axis, and the jet is uniform or non-uniform(the energy per unit solid angle decreases smoothly away from the axis $\\epsilon (\\theta)\\propto (\\theta/\\theta_c)^{-k}$). We find that, for uniform jet($k=0$), the afterglow light curves for different viewing angles are somewhat different: in general, there are two breaks in the light curve, the first one corresponds to the time at which $\\gamma\\sim (\\theta_j-\\theta_v)^{-1}$, and the second one corresponds to the time when $\\gamma\\sim (\\theta_j+\\theta_v)^{-1}$. However, for non-uniform jet, the things become more complicated. For the case $\\theta_v=0$, we can obtain the analytical results, for $k<8/(p+4)$(where $p$ is the spectral index of electron energy distribution) there should be two breaks in the light curve correspond to $\\gamma\\sim\\theta_c^{-1}$ and $\\gamma\\sim\\theta_j^{-1}$ respectively, while for $k>8/(p+4)$ there should be only one break corresponds to $\\gamma\\sim\\theta_c^{-1}$, and this provides a possible explanation for some rapidly fading afterglows whose light curves have no breaks since the time at which $\\gamma\\sim\\theta_c^{-1}$ is much earlier than our first observation time. For the case $\\theta_v\\neq 0$, our numerical results show that, the afterglow light curves are strongly affected by the values of $\\theta_v$, $\\theta_c$ and $k$. If $\\theta_v$ is close to $\\theta_c$ and $k$ is small, then the light curve is similar to the case of $k=0$, except the flux is somewhat lower. However, if the values of $\\theta_v/\\theta_c$ and $k$ are larger, there will be a prominent flattening in the afterglow light curve, which is quite different from the uniform jet, and after the flattening a very sharp break will be occurred at the time $\\gamma\\sim (\\theta_v + \\theta_c)^{-1}$. ",
        "introduction": "Gamma-ray bursts (GRBs) are known as an explosive phenomenon occurring at cosmological distances, emitting large amount of energy mostly in the gamma-ray range (see, e.g.  Piran 1999; Cheng \\& Lu 2001 for a review). Observations show that some of GRBs are emitting an extremely large energy with $E_\\gamma\\gg 10^{52}$ ergs if emission is isotropic. For example, GRB990123, the most energetic GRB event detected so far, has an isotropic gamma-ray energy of $\\sim 3.4\\times 10^{54}$ ergs, which corresponds to the rest-mass energy of $\\sim 1.9M_{\\odot}$ (Kulkarni et al. 1999). Such a crisis of extreme large energy forced some people to think that the GRB emission must be highly collimated in order to reduce the total energy.  The second reason for GRB emission being jet-like comes from the fact that there are sharp breaks in the light curves of some GRBs' afterglows, such as GRB990123 (Kulkarni et al. 1999; Castro-Tirado et al. 1999) and GRB990510 (Harrison et al. 1999; Stanek et al. 1999), etc.. These observed breaks have generally been interpreted as evidence for collimation of the GRB ejecta, since Rhoads (1999) and Sari et al. (1999) have pointed out that the lateral expansion of the relativistic jet can produce a sharp break in the afterglow light curve.  However, in the current afterglow jet models, it is generally assumed that the jet is uniform, and the line-of-sight is just along the jet axis. It is obvious that these assumptions are usually not true, since some GRB models predict that the jet may be non-uniform, within which the energy per unit solid angle decreases away from the jet axis, such as $\\epsilon\\propto\\theta^{-k}$ (e.g. MacFadyen et al. 2001), and the probability that the line-of-sight just crosses the jet axis is near zero. So it is very important to investigate the situation that the jet is non-uniform and is seen by observers at a wide range of viewing angles from the jet axis. Some authors have already considered the situation of anisotropic jets (Meszaros, Rees \\& Wijers 1998; Salmonson 2001; Dai \\& Gou 2001; Zhang \\& Meszaros 2002; Rossi, Lazzati \\& Rees 2002). In this paper, we give a detailed calculation of the emission from anisotropic jets, including the effect of equal-arrival-time surface. In the next section we discuss the dynamical evolution of the jet, in section 3 we calculate the emission from uniform jet for different viewing angle, in section 4 we calculate the emission features from non-uniform jet, and finally we present some discussions and conclusions. ",
        "conclusions": "In this paper we calculate the GRB afterglow light curves from relativistic jets in more details, assuming that the jet may be uniform or non-uniform, and the observer locate at arbitrary angle with respect to the jet axis. We have shown that there are several distinct features of our jet emission compared with previous jet model.  In previous analysis, it is generally assumed that the jet is uniform and the line-of-sight is just along the jet axis, in this case the afterglow light curves have a break at the time $\\gamma\\sim \\theta_j^{-1}$. However, if the viewing angle $\\theta_v \\neq 0$, we have shown that there should be two breaks in the light curve, the first one corresponds to the time at which $\\gamma\\sim (\\theta_j-\\theta_v)^{-1}$, and the second one corresponds to the time when $\\gamma\\sim (\\theta_j+\\theta_v)^{-1}$, although these transitions are very smoothly.  If the jet is not uniform, within which the energy distribution is given by eq.(10), then the calculation is more complicated. In the case of $\\theta_v =0$, we can give an analytical result, for $k<8/(p+4)$(where $p$ is the spectral index of electron energy distribution) there should be two breaks in the light curve correspond to $\\gamma\\sim\\theta_c^{-1}$ and $\\gamma\\sim\\theta_j^{-1}$ respectively, while for $k>8/(p+4)$ there should be only one break corresponds to $\\gamma\\sim\\theta_c^{-1}$, after that the flux decays as $F_\\nu\\propto T^{-3p/4}$. We argue that this may explain some rapidly fading afterglows whose light curves have no breaks, since the time $T_1$, at which $\\gamma\\sim\\theta_c^{-1}$, is usually earlier than our first observation time.  If the jet is non-uniform and the viewing angle $\\theta_v\\neq 0$, only numerical results can be given. We have shown that in this case the shape of afterglow light curve is dependent on the values of $\\theta_v$, $\\theta_c$ and $k$. When $\\theta_v$ is near $\\theta_c$ and $k$ is small, then the light curve is similar to the case of $\\theta_v=0$, except the flux is somewhat lower. However, if the viewing angle $\\theta_v$ is larger than $\\theta_c$ and $k$ is not small, then there will be a prominent flattening in the afterglow light curve, which is quite different from the case of uniform jet, and after the flattening a very sharp break will be occurred around the time $\\gamma\\sim (\\theta_v + \\theta_c)^{-1}$. We think this is a main difference between the uniform and non-uniform jet, and we can identify whether the jet is uniform or not by this feature.  It is not very difficult to understand why sometimes there is a flattening in the light curve. It is well known that, for a relativistic blast wave with Lorentz factor $\\Gamma\\gg 1$, the observer can only observe a solid angle around $\\theta_v$ with a half opening angle of order $\\Gamma ^{-1}$, the contribution from other components can be neglected, so when the blast wave decelerates, the observer can see larger components. For the case of non-uniform jet and $\\theta_v >\\theta_c$, the energy at $\\theta=\\theta_v$ is much smaller than that of smaller $\\theta$, so when $\\Gamma$ decreases, the observer can observe more region with larger energy (smaller $\\theta$), so there will be a flattening in the light curve.  We would point out that, in our calculation we have neglected the sideways expansion of the jet since this process is too complicated. However, we know that, in fact this process is an important issue in determining the shape of a light curve, since it will likely significantly change the shape of the light curve for both the uniform and non-uniform jets. So we suggest that for a more realistic calculation the sideways expansion should be taken into account."
    },
    "0212/astro-ph0212208_arXiv.txt": {
        "abstract": "The off-axis location of the Advanced Camera for Surveys (ACS) is the chief (but not sole) cause of strong geometric distortion in all detectors: the Wide Field Camera (WFC), High Resolution Camera (HRC), and Solar Blind Camera (SBC).  Dithered observations of rich star cluster fields are used to calibrate the distortion.  We describe the observations obtained, the algorithms used to perform the calibrations and the accuracy achieved. ",
        "introduction": "Images from the Hubble Space Telescope (HST) Advanced Camera for Surveys (ACS) suffer from strong geometric distortion: the square pixels of its detectors project to trapezoids of varying area across the field of view.  The tilted focal surface with respect to the chief ray is the primary source of distortion of all three ACS detectors.  In addition, The HST Optical Telescope Assembly induces distortion as does the ACS M2 and IM2 mirrors (which are designed to remove HST's spherical aberration).  The SBC's optics include a photo-cathode and micro-channel plate which also induce distortion.  Here we describe our method of calibrating the geometric distortion using dithered observations of star clusters.  The distortion solutions we derived are given in the IDC tables delivered in Nov 2002, and currently implemented in the STScI CALACS pipeline.  This paper is a more up to date summary of our results than that presented at the workshop.  An expanded description of our procedures is given by Meurer (2002). ",
        "conclusions": ""
    },
    "0212/astro-ph0212147.txt": {
        "abstract": "We present rest-frame $R$-band galaxy luminosity function measurements for three different redshift ranges: $0.5\\leq z \\leq 0.75$, $0.75 \\leq z \\leq 1.0$, and $1.0\\leq z \\leq 1.5$.  Our measurements are based on photometric redshifts for $\\sim 3000$ $H$-band selected galaxies with apparent magnitudes $17\\leq H \\leq 20$ from the Las Campanas Infrared Survey.  We show that our photometric redshifts are accurate with an RMS dispersion between the photometric and spectroscopic redshifts of $\\sigma_z/(1+z) \\approx 0.08$.  Using galaxies identified in the Hubble Deep Field South and Chandra Deep Field South regions, we find, respectively, that ($7.3\\pm 0.2$)\\,\\% and ($16.7\\pm 0.4$)\\,\\% of the $H\\leq 20$ galaxies are at $z\\geq 1$.  We first demonstrate that the systematic uncertainty inherent in the luminosity function measurements due to uncertainties in photometric redshifts is non-negligible and therefore must be accounted for.  We then develop a technique to correct for this systematic error by incorporating the redshift error functions of individual galaxies in the luminosity function analysis.  The redshift error functions account for the non-gaussian characteristics of photometric redshift uncertainties.  They are the products of a convolution between the corresponding redshift likelihood functions of individual galaxies and a Gaussian distribution function that characterizes template-mismatch variance.  We demonstrate, based on a Monte Carlo simulation, that we are able to completely recover the bright end of the intrinsic galaxy luminosity function using this technique.  Finally, we calculate the luminosity function separately for the total $H$-band selected sample and for a sub-sample of early-type galaxies that have a best-fit spectral type of E/S0 or Sab from the photometric redshift analysis.  The primary results of this analysis are: (1) the galaxy luminosity functions are consistent with a Schechter (1976) form; (2) the evolution of the co-moving luminosity density $\\ell$ of $H$-band selected galaxies is characterized by $\\Delta\\log \\ell\\,/\\Delta\\log (1+z)=0.6 \\pm 0.1$ at rest-frame 6800\\,\\AA; and (3) $\\ell_R$ for color-selected early-type galaxies exhibits moderate evolution from $z\\sim 1.5$ to $z \\sim 0.3$.  Specifically, $\\ell_R$ for these red galaxies brighter than 1.0 (1.6) $L_*$ could decrease with increasing redshift by {\\em at most} a factor of three (six) over this redshift range after removing possible stellar fading according to the most extreme stellar evolution scenario. ",
        "introduction": "The galaxy luminosity function represents the product of a series of physical processes that govern galaxy evolution, including star formation and the recycling of metals through various heating and cooling processes.  An accurate estimate of the galaxy luminosity function as a function of time therefore bears significantly on our understanding of how structures in the universe form and evolve.  Independent measurements of the optical galaxy luminosity function based on various deep redshift surveys between $z\\sim 0$ and $z\\sim 0.75$ have yielded consistent results (Lilly \\etal\\ 1995; Ellis \\etal\\ 1996; Marzke \\& da Costa 1997; Bromley \\etal\\ 1998; Marzke \\etal\\ 1998; Lin \\etal\\ 1999; Madgwick \\etal\\ 2001).  For example, the galaxy luminosity function is well represented by a Schechter function (Schechter 1976) with a steeper faint-end slope for galaxies of bluer optical colors at all epochs that have been studied. Comparisons of these measurements at different redshifts further show that these blue galaxies exhibit a stronger evolution than red galaxies both in their space density and in their rest-frame optical luminosity.  It is difficult, however, to apply these results to fully constrain galaxy formation models.  While galaxy luminosities measured in the optical are sensitive to the transient star formation rate and therefore are a good tracer of the instantaneous star formation activity, they are not representative of the total stellar mass that has been assembled in the galaxies over time. Additional caveats exist in all studies that attempt to draw conclusions based on comparisons of luminosity function measurements obtained using different galaxy samples.  For example, comparison of morphologically-selected and color-selected galaxy samples is inadequate, because even morphologically distinct galaxies can exhibit similar colors.  Furthermore, it is impossible to unambiguously identify the same galaxy population at different epochs using the observed colors because the underlying stellar evolution sequence is unknown. Finally, other selection biases such as the cosmic surface brightness dimming effect (preventing accurate morphological classifications at high redshift), bandpass selection effect (galaxies selected in the same observed-frame bandpass may have different rest-frame properties), and the presence of dust (in which a young, star-forming galaxy masquerades as an old, quiescent one) further complicate the interpretations of comparisons of different galaxy surveys.  We have initiated the Las Campanas Infrared (LCIR) survey (Marzke \\etal\\ 1999; McCarthy \\etal\\ 2001; Chen \\etal\\ 2002; Firth \\etal\\ 2002) to study the history of mass assembly based on a uniform sample of near-infrared selected galaxies.  The rest-frame near-infrared light is a good proxy for the total stellar mass, independent of age or galaxy type (Charlot 1996).  Measurements of the galaxy luminosity function at rest-frame near-infrared wavelengths at different epochs may, therefore, be adopted to estimate the stellar mass density evolution.  A robust measurement of the rest-frame near-infrared luminosity function at $z \\approx 1$ requires a galaxy survey conducted in the $K$ band.  We have so far completed the LCIR $H$-band survey and a complete $K$-band survey is underway.  Here we present measurements of the rest-frame $R$-band luminosity functions in three redshift intervals over the range $z=0.5$--1.5 for the $H$-band selected galaxies using photometric redshifts. At $z\\approx 1.2$, the observed-frame $H$-band approximately corresponds to rest-frame $R$-band, and therefore the calculation does not depend sensitively on the templates adopted to determine the $k$-correction.  The primary objectives of the analysis are (1) to obtain measurements of the galaxy luminosity function at $z\\geq 1$, where different galaxy formation scenarios begin to show distinct predictions in the space densities and masses of evolved galaxies, and (2) to study galaxy evolution at redshifts between $z\\approx 0.5$ and $z\\approx 1.5$ based on a uniform sample of near-infrared selected galaxies, thus extending the redshift range of existing deep redshift surveys to beyond $z=1$.  We first conducted a photometric redshift survey in the LCIRS galaxy sample to identify a complete sample of galaxies at $z>0.5$.  The photometric redshift survey in two of the survey fields identified $\\geq 3000$ galaxies at $z\\geq 0.5$, yielding a statistically significant sample of intermediate-redshift galaxies for understanding the galaxy population at redshifts beyond $z=1$. The sample size is comparable to those of wide-field spectroscopic surveys carried out at $z\\leq 0.5$ such as the Canadian Network for Observational Cosmology Field Galaxy Redshift Survey (CNOC2; e.g.\\ Lin \\etal\\ 1999).  The photometric redshift techniques determine galaxy redshifts based on the presence/absence of broad-band spectral discontinuities, rather than narrow-band emission or absorption line features.  They are therefore of particular value for the identification of galaxies at $z\\geq 0.75$.  Complete {\\em spectroscopic} surveys of faint galaxies in this redshift range are extremely challenging, as prominent spectral line features in galaxy spectra are redshifted to near-infrared wavelengths and are difficult to detect due to bright, contaminating sky lines and detector limitations.  It has been demonstrated over the past several years that distant galaxies may be robustly identified using photometric redshift techniques that incorporate optical and near-infrared broad-band photometry (Connolly \\etal\\ 1997; Lanzetta, Fern\\'{a}ndez-Soto, \\& Yahil 1998; Fern\\'{a}ndez-Soto, Lanzetta, \\& Yahil 1999; Martini 2001; Fern\\'{a}ndez-Soto \\etal\\ 2001; Rudnick \\etal\\ 2001).  In this paper, we show that the photometric redshift measurements of the galaxies identified in the LCIR survey are accurate, with an RMS dispersion between the photometric and spectroscopic redshifts of $\\sigma_z/(1+z) \\approx 0.08$.  Next, we developed a maximum likelihood method to measure galaxy luminosity function using photometric redshifts, which explicitly accounts for photometric redshift uncertainties of individual galaxies.  Photometric redshift uncertainties together with a steep slope of the galaxy luminosity function would flatten the observed luminosity function and yield systematic biases. Uncertainties in photometric redshifts must therefore be accounted for in the luminosity function analysis.  We demonstrate that systematic uncertainties in the derived galaxy luminosity function are significantly reduced when redshift error functions are taken into account.  We show that (1) the galaxy luminosity function is consistent with a Schechter (1976) form; (2) the evolution of the co-moving luminosity density of the $H$-band selected galaxy is characterized by $\\Delta \\log \\ell\\,/\\Delta\\log (1+z)=0.6 \\pm 0.1$ at rest-frame 6800\\,\\AA; (3) galaxies with spectral energy distributions best characterized as E/S0 or Sab appear to have faded by $\\approx 0.5$ mag from $z \\sim 1$ to $z\\sim 0.5$, consistent with the expected luminosity evolution from an exponentially declining star formation rate model; and (4) only moderate evolution in the space density and in the rest-frame, co-moving $R$-band luminosity density is found for the bright, color-selected early-type galaxies.  We adopt the following cosmology: $\\Omega_{\\rm M}=0.3$ and $\\Omega_\\Lambda= 0.7$ with a dimensionless Hubble constant $h = H_0/(100 \\ {\\rm km} \\ {\\rm s}^{-1}\\ {\\rm Mpc}^{-1})$ throughout the paper. ",
        "conclusions": "\\renewcommand{\\thefootnote}{\\fnsymbol{footnote}} \\setcounter{footnote}{0}  We have developed a technique to obtain a robust estimate of the intrinsic luminosity function using photometric redshifts.  We adopt an ``empirical'' estimate of the redshift error function for each galaxy, which is a convolution of the redshift likelihood function that characterizes the uncertainty due to photometric errors and a Gaussian kernel of 1-$\\sigma$ width $\\sigma_z$ that characterizes the uncertainty due to template mismatch.  This technique takes into account the non-gaussian error characteristics of photometric redshifts, and thus differs from the approach of Subbarao \\etal\\ (1996), who used a simple Gaussian to model the photometric redshift uncertainties.  The results of the simulation described in \\S\\ 3.4 show that our approach allows us to completely recover the galaxy luminosity function to the sensitivity limit of the survey. Here we compare the derived galaxy luminosity functions at different redshifts and discuss the implications for the evolution of galaxies since $z\\sim 1.5$.  \\subsection{Galaxy Evolution at $0.5\\leq z \\leq 1.5$}  It is clear from Figure 9 that the depth of the $H$-band survey is not sufficiently sensitive to constrain the statistical properties of galaxies fainter than $\\sim L_*$ at $z\\geq 0.75$, and we are therefore unable to determine the evolution of the faint galaxy population based on the $H$-band survey.  However, the data are sufficient to determine the evolution of the bright galaxy population at $0.5\\leq z \\leq 1.5$ based on the $H$-band survey data.  The evolution of this population has significant implications for understanding how galaxies form because different formation scenarios make distinct predictions for the number density evolution of massive galaxies.  For example, under monolithic collapse scenarios (e.g.\\ Eggen, Lynden-Bell, \\& Sandage 1962), massive galaxies form early on a dynamical timescale and have a constant co-moving space density and gradually declining intrinsic luminosities.  In contrast, hierarchical formation scenarios (e.g.\\ White \\& Rees 1978) predict that massive galaxies form through the merger of lower-mass galaxies over a Hubble time and therefore the co-moving space density of massive galaxies increases as the universe evolves.  The luminosity functions presented in Figure 9 and Tables 1 and 2 show several interesting results.  First, the best-fit Schechter luminosity functions agree with the stepwise luminosity functions to within the measurement uncertainties at all redshifts, including the luminosity function at $0.5\\leq z\\leq 0.75$ where our survey is sufficiently sensitive to identify galaxies as faint as $\\sim 0.2\\,L_*$.  Second, there exists little or no evolution either in $M_{R_*}$ or in $\\phi_*$ for the total $H$-band selected population at $0.5\\leq z \\leq 1.5$.  Third, we find a moderate luminosity evolution for galaxies with an SED best characterized as E/S0 or Sab galaxies. Specifically, our measurements obtained from the modified STY approach suggest that from $z\\sim 1$ to $z\\sim 0.5$ these early-type galaxies have faded by $\\sim 0.5$ mag.  This is consistent with the expected luminosity evolution for a galaxy formed at $z_f=30$ and following an exponentially declining star formation rate model with $\\tau = 1$ Gyr (McCarthy \\etal\\ 2001; Chen \\etal\\ 2002).  Finally, we find only a mild evolution in $\\phi_*$ between $z\\approx 0.5$ and $z \\approx 1$ ($\\sim 40$\\% decrease with redshift) for the color-selected early-type galaxies, but we cannot argue against the no evolution scenario at more than the $2\\,\\sigma$ level.  When we consider the entire $H$-band selected sample at $0.5\\leq z \\leq 1.5$, our estimates of $\\phi_*$ are completely consistent with a no evolution scenario.  It appears that neither the total $H$-band selected sample nor the color-selected early-type galaxies have evolved significantly in intrinsic luminosity or space density since $z\\sim 1.5$, contrary to the predictions of the hierarchical formation scenarios.  But as with all luminosity function measurements, it is clear that there exists a strong degeneracy between $M_{R_*}$ and $\\phi_*$ at $z\\geq 0.75$ in our analysis, because the $H$-band survey is not sensitive to sub-$L_*$ galaxies at this redshift range and our estimates for these two parameters were determined from the space density of intrinsically bright galaxies with a fixed faint-end slope $\\alpha$.  For example, a bright $M_{R_*}$ together with a small $\\phi_*$ can produce the same bright-end luminosity function as a faint $M_{R_*}$ with a large $\\phi_*$.  To estimate how significantly early-type galaxies have evolved since $z\\sim 1.5$, we therefore calculated the rest-frame co-moving $R$-band luminosity density $\\ell_R$ for galaxies brighter than $M_{\\rm min}$ using the best-fit Schechter luminosity function, \\begin{eqnarray} \\ell_R &=& \\int_{L_{min}}^\\infty L\\cdot \\Phi(L;L_*)\\,d(L/L_*) \\\\ &=& \\int_{-\\infty}^{M_{\\rm min}} (0.4 \\ln 10)\\cdot\\phi_*\\cdot L_*\\cdot 10^{0.4\\,(M_*-M)(2+\\alpha)}\\cdot\\exp(-10^{0.4\\,(M_*-M)})\\,dM. \\end{eqnarray} This represents an integrated quantity that characterizes the brightest galaxy population and is least sensitive to faint galaxies.  We test different formation scenarios using the estimated co-moving luminosity density evolution of the bright, red galaxies.  Including the luminosity function estimated for early-type galaxies at $z=0.3$ in the CNOC2 survey (Lin \\etal\\ 1999), we first compared $\\ell_R$ calculated for a constant $M_{\\rm min}$ at all redshifts.  The results are presented in the left panel of Figure 11, which shows a flat $\\ell_R$ versus $z$ for $M_{\\rm min} = -20.5$ (squares) , $-21.0$ (triangles), and $-21.5$ (circles).  These roughly corresponds to $L_*$, $1.6\\,L_*$, and $2.5\\,L_*$, respectively, at $z\\sim 0.3$. However, because stars evolve with time, the corresponding luminosity evolution must be accounted for.  To distinguish between a passively evolving population and hierarchical mergers, we estimated the amount of brightening ($\\Delta\\, M_R$) with increasing redshift (or with decreasing age of the universe) for different stellar evolution models (see Table 3 for a summary).   We considered a wide range of stellar evolution scenarios, from a single burst at $z_f=30$, to an exponentially declining star formation rate model with $z_f=5$ and $\\tau = 1$ Gyr.  The former represents a classic passive evolution mode under the monolithic collapse scenariol.  The latter represents the closest resemblance of a more continuous star-forming scenario in which some residual star formation is taking place at $z<3$, while there is still enough time for the red galaxy population to be present by $z\\sim 1$.  We then calculated $\\ell_R$ for a suite of evolving $M_{\\rm min}(z)$ according to different stellar evolution recipes.  The goal of this analysis is to examine whether or not mergers are the dominant process for the formation of red galaxies.  Mergers would lead to brightening with time (smaller systems combined to form bigger ones), which compensates the expected fading due to stellar evolution.  We would therefore expect to observe an increasing $\\ell_R$ with decreasing redshift using ($\\Delta\\,M_R$) presented in Table 3, if these red galaxies formed through a sequence of mergers.  A flat redshift distribution of $\\ell_R$ would then imply that mergers are insignificant in the redshift interval tested.  We present the results in Figure 11.  The middle panel is for the $z_f=30$ single burst scenario and the right panel is for the $z_f=5$ exponentially declining star formation rate model.  The results in these two panels have been scaled to have consistent $M_{\\rm min}$ at $z=0.3$ as in the left panel.  In addition, we exclude the brightest subsample (circles in the left panel) from these two panels, because the number of galaxies at the very bright end of the luminosity functions is very small beyond $z\\sim 0.75$.  Specifically, there are only eight galaxies in the red sample that are brighter than $M_R=-23$ at $1.0 \\leq z \\leq 1.5$.  Figure 11 shows that $\\ell_R$ evolves relatively slowly with redshift for different bright subsamples, and that it could have increased by at most a factor of $\\approx 6$ from $z\\sim 1.2$ to $z\\sim 0.3$ for the brightest early-type sample under the most extreme stellar evolution model presented in the right panel.  Taking into account the fact that these models do not have dust included, which may further reduce the amount of fading in intrinsic luminosity with redshift, we find no evidence to support that mergers dominate the formation of early-type galaxies over $0.5 \\leq z \\leq 1.5$.  This is consistent with the findings of Firth \\etal\\ (2002), who directly compared the observed abundances of red galaxies at $z\\sim 1$ with those from different semi-analytical models.  They found that the models underpredict the number of red galaxies observed in the LCIR survey by at least a factor of three.  The rest-frame co-moving $R$-band luminosity densities for all $H$-band selected galaxies at different redshifts are presented in Figure 12.  As in Figure 11, we calculated $\\ell_R$ for three constant thresholds, $M_{\\rm min}= -20.5$ (squares) , $-21.0$ (triangles), and $-21.5$ (circles).  Because galaxies in the total $H$-band selected sample have a wide range of star formation history, we cannot remove the effects of pure luminosity evolution by adopting a simple $M_{\\rm min}(z)$.   The results of Figure 12 show that there exists only a mild evolution in the rest-frame co-moving $R$-band luminosity density for the entire galaxy population over $0.5 \\leq z\\leq 1.5$.  It is also important to note that our analysis is based on an $H$-band selected galaxy sample, which roughly corresponds to galaxies selected in the rest-frame $J$, $I$, and $R$ bands at $z\\sim 0.5$, 0.9, and 1.3, respectively. A potential systematic bias in our measurements is clearly the bandpass selection effect which, when combined with variation in star formation rate and stellar evolution with time, makes it difficult to interpret our results, particularly for the total $H$-band selected sample.  At $z\\sim 1.3$, the $H$-band survey preferentially selects optically bright (and therefore younger) galaxies, while it tends to select near-infrared bright (and therefore older) galaxies at $z\\sim 0.5$, when the universe was twice as old.  Because the total $H$-band selected sample consists of galaxies of different star formation histories in different redshift intervals, it is possible that galaxies of different colors, selected at different rest-frame wavelengths at different redshifts conspire to produce a luminosity function that appears to show little evolution.  Consequently, it is impossible to rule out any evolution scenarios without a detailed understanding of the intrinsic properties of the galaxies in the total $H$-band sample.  On the other hand, the bandpass selection bias has a much less serious effect for the early-type galaxy sample. Only galaxies with an observed SED (spanning a spectral range from the $V$ band or $U$ for the HDFS objects through the $H$ band) that is consistent or redder than a typical Sab galaxy were included in the analysis at all redshifts.  Our measurements represent a true, empirical measure of how ``red'' galaxies evolve with time, and the results of our analysis suggest that early-type galaxies have not evolved significantly since $z\\sim 1.5$.  \\subsection{Comparison with Previous Studies}  The luminosity functions presented in \\S\\ 3.6 cover the redshift range $0.5 \\leq z \\leq 1.5$.  The upper end of this range extends beyond the reach of traditional spectroscopic redshift surveys, the most ambitious of which stretch to $z=1.2$ but become significantly incomplete beyond $z=1$ (Lilly \\etal\\ 1995; Cowie \\etal\\ 1996; Cohen 2002).  Various groups have measured evolution in the luminosity function using spectroscopic redshifts over the redshift range $0 < z < 1$ for galaxies of different colors (e.g.\\ Lilly \\etal\\ 1995), different spectral types (e.g.\\ Heyl \\etal\\ 1997; Cohen 2002), or different morphology (Im \\etal\\ 1999).  In addition, some groups have attempted to measure the luminosity function using photometric redshifts for galaxies at $z\\apl 4$ (Gwyn \\& Hartwick 1996; Mobasher \\etal\\ 1996; Connolly \\etal\\ 1997; Sawicki \\etal\\ 1997).  Different sample selection criteria make a quantitative comparison between these measurements difficult.  Here we present a qualitative comparison between our measurements and the previous ones.  First, we consider previous measurements that were conducted in the rest-frame $R$ band.  The {\\it local} luminosity density is now well constrained by both the 2dF Redshift Survey (Cross \\etal\\ 2001) and the SDSS (Blanton 2001).  It is encouraging that these surveys yielded consistent results in the $B$ band.  The $R$-band luminosity density available from the SDSS is most directly relevant to our results from the $H$-selected LCIR survey, and is included in Figure 12 for the luminosity density evolution of the total $H$-band selected sample.  If we consider only the LCIR measurements at $0.5 < z < 1$ and the SDSS measurement at $z < 0.2$, then the evolution in the luminosity density from $z \\sim 0$ to $z\\sim 1$ is less than 0.2 dex for any reasonable choice of the faint-end slope, $\\alpha$, indicating a very slow evolution.  Next, we compare our results with those from a number of very deep redshift surveys that covered a relatively smaller area but were sensitive to galaxies at redshift beyond $z\\sim 1$.  These are the Canada-France Redshift Survey (CFRS; e.g.\\ Lilly \\etal\\ 1996), a deep $K$-band survey by Cowie, Songaila, \\& Barger (1999), and the Caltech Faint Galaxy Redshift Survey (CFGRS; Cohen \\etal\\ 2002), all of which showed that the bulk of the evolution in the luminosity function occured in bluer spectral types.  In addition, they showed that galaxies with redder SEDs evolved more slowly, if at all, over this redshift range, a result with which the LCIR survey is at least qualitatively consistent.  Note however that these results may change if adopting a different color selection criterion for red galaxies at high redshifts (Kauffmann, Charlot, \\& White 1996) or the incompleteness correction for absorption feature dominated galaxies was incorrect (Totani \\& Yoshii 1997), two effects that work in opposite ways.  A more quantitative comparison is possible with Lilly \\etal\\ (1996), which gave the evolution of the luminosity density between $z=0$ and $z=1$ for different CFRS samples.  In particular, Figure 12 is most directly comparable to Figure 1 in Lilly \\etal\\ (1996).  The CFRS luminosity densities were computed at rest-frame 2800\\,\\AA, 4400\\,\\AA, and 1\\,$\\mu$m by interpolating their $BVIK$ photometry, which showed a decline by approximately 0.5 dex at 4400\\,\\AA\\ and 0.4 dex at $1\\,\\mu$m from $z\\sim 1$ to $z\\sim 0$.\\footnote{We have subtracted 0.2 dex from the original measurements in Lilly \\etal\\ to account for the differences in the adopted cosmological models.}  This gives $\\Delta\\log \\ell\\,/\\Delta\\log (1+z) = 1.7 \\pm 0.5$ at 4400\\,\\AA\\ and $1.3 \\pm 0.4$ at $1\\,\\mu$m.  We find from the LCIR survey that $\\Delta\\log \\ell\\,/\\Delta \\log (1+z)=0.6 \\pm 0.1$ at 6800\\,\\AA, if we assume a faint-end slope $\\alpha= -1.0$.  If we assume a very steep faint-end slope $\\alpha=-1.5$, then we measure still only $1.0 \\pm 0.2$.  Although our results are marginally consistent with the CFRS results, the $H$-selected LCIR survey appears to favor slower evolution in the overall galaxy population at $z < 1$ (at the $\\approx \\,1\\,\\sigma$ level).  The slow evolution measured here is similar to the $I$-band evolution measured by Cowie, Songaila, \\& Barger (1999) from $K$-selected spectroscopic surveys at $z < 1.2$.  Based on a much larger sample covering a broader redshift range, we find that this slower evolution at red wavelengths extends to at least $z=1.5$.  Our results are also consistent with the interpretation of the luminosity function evolution derived from the CFGRS, which represents by far the largest and most complete spectroscopic sample of galaxies at $z > 0.8$.  It is, however, still limited by small sample size (144 galaxies at $0.8 \\leq z \\leq 1.05$ and 18 galaxies at $1.05 \\leq z \\leq 1.5$).  This illustrates the difficulties of carrying out a complete spectroscopic survey over a large area to very faint magnitude limits.  Next, we compare our results with those estimated based on photometric redshifts.  At $1 < z < 2$, photometric redshift surveys prevail but have so far been limited to small areas of the sky.  In addition, it is important to include near-infrared photometry in order for accurate photometric redshift measurements to be plausible.  Connolly \\etal\\ (1997) showed that uncertainties of photometric redshifts are reduced by 40\\%, when incorporating near-infrared photometry in the photometric redshift analysis.  Indeed, the reddest galaxies known in this reshift interval (Elston \\etal\\ 1988; McCarthy \\etal\\ 1992; Hu \\& Ridgeway 1994) are often exceedingly faint in all optical bandpasses bluer than the $R$ band, making photometric redshift measurements very uncertain.  The luminosity function from Connolly \\etal\\ was for a $J$-band selected galaxy sample over $1 < z < 2$ in the Hubble Deep Field North (HDFN).  Their luminosity function was computed, however, at rest-frame 2800 \\AA, which is more sensitive to young stellar populations than the rest-frame $R$ band. Furthermore, over the redshift range probed by the LCIR survey ($0.5 < z < 1.5$), the HDFN is primarily sensitive to the faint end of the luminosity function, while the LCIR survey is sensitive only to the bright end.  Finally, the total survey area of HDFN was only 4.48 arcmin$^2$, corresponding to a comoving volume of $\\approx 1.3\\times 10^4$ Mpc$^3$ over the $1 \\leq z \\leq 2$ interval.  In contrast, our analysis is based on a survey volume that is more than two orders of magnitude larger than in the HDFN.  Therefore, we have sufficient numbers of bright galaxies to $z=1.5$ to probe their evolution in greater detail.  The larger volume is necessary to accurately measure the mean space density because the bright, red galaxies are now known to be strongly clustered (Daddi \\etal\\ 2000; McCarthy \\etal\\ 2001; Firth \\etal\\ 2002).  Other groups have attempted to study the evolution of the galaxy luminosity function at $z\\leq 1$ based on a deep photometric redshift survey that incorporates near-infrared photometry (Fried \\etal\\ 2001; Poli \\etal\\ 2001), but these galaxy samples were selected in the $I$ band, which lends a bluer tint to the overall mix of SEDs than the $H$-selected LCIR survey galaxies.  It is therefore not straightforward to compare our results with these other photometric redshift surveys.  Finally, we consider measurements obtained for early-type galaxies selected on the basis of quantitative morphological criteria such as smoothness, symmetry, and bulge-to-total ratio.  Im \\etal\\ (1999) selected a sample of 145 E/S0 galaxies from deep HST imaging in the Groth Strip (Groth \\etal\\ 1994) Using a subset with spectroscopic redshifts as calibrators, they estimated redshifts for the remaining galaxies using the $V-I$ color alone.  The resulting luminosity functions (measured in the rest-frame $B$ band) suggest slow evolution in the E/S0 population at $z < 1$, in qualitative agreement with the results discussed earlier.  Brinchmann \\etal\\ (1998) and Menanteau \\etal\\ (1999) have, however,suggested that the bulk of morphologically selected early-type galaxies at $z < 1$ have bluer colors than those predicted by different stellar evolution models, but the results are inconclusive because they depend sensitively on the adopted models and how photometric uncertainties are incorporated in these models.  In summary, we find that the rest-frame, co-moving optical luminosity density of the entire galaxy population has evolved only moderately (declining by a factor of $\\approx 2$) from $z\\sim 1$ to $z\\sim 0$.  The rest-frame, co-moving optical luminosity density of the color-selected, early-type galaxies has not evolved significantly since $z\\sim 1$.  Our new measurements of luminosity density evolution is based on a significantly larger sample selected in the near-infrared than previous surveys selected at a range of wavelengths.  It is consistent with previous findings, but we place a much tighter constraint on the apparant lack of evolution in the co-moving luminosity density of early-type galaxies."
    },
    "0212/astro-ph0212452_arXiv.txt": {
        "abstract": "Gamma-ray burst detectors are sensitive at different energies, complicating the comparison of the burst populations that they detect.  The instrument teams often report their detector sensitivities in their instruments' energy band.  I propose that sensitivities be reported as the threshold peak photon flux $F_T$ over the 1--1000~keV energy band for a specific spectral shape. The primary spectral parameter is $E_p$, the energy of the maximum $E^2 N_E\\propto \\nu f_\\nu$.  Thus $F_T$ vs. $E_p$ is a useful description of a detector's sensitivity.  I find that Swift will be marginally more sensitive than BATSE for $E_p>100$~keV, but significantly more sensitive for $E_p<100$~keV.  Because of its low energy sensitivity, the FREGATE on {\\it HETE-2} is surprisingly sensitive below $E_p=100$~keV. Both the WFC on {\\it BeppoSAX} and the WXM on {\\it HETE-2} are/were sensitive for low $E_p$. As expected, the GBM on {\\it GLAST} will be less sensitive than BATSE, while {\\it EXIST} will be significantly more sensitive than Swift.  The {\\it BeppoSAX} GRBM was less sensitive that the WFC, particularly at low $E_p$. ",
        "introduction": "The gamma-ray burst missions that have flown in the past 15 years have reported that bursts are hard (e.g., BATSE---Preece et al. 2000; Mallozzi et al. 1995) or soft (e.g., Ginga---Strohmayer et al. 1998).  However these missions had different sensitivities to hard and soft bursts, and have reported the threshold burst intensities using a variety of different measures, such as peak flux or energy fluence measured over different energy bands. To synthesize the results of these mission (and determine whether they are mutually consistent), and to compare the capabilities of past, current, and proposed missions, we need common measures of burst intensity and hardness, and we need to express a burst detector's capabilities in terms of those common measures.  I propose to characterize bursts by the peak photon flux integrated over 1--1000~keV averaged over one second, and by the spectrum's peak energy $E_p$ during this one second. The peak energy $E_p$ is the energy of the maximum of $N(E)E^2$ (proportional to $\\nu F_\\nu$)---the energy flux per logarithmic energy (frequency) band; $E_p$ is a first order measure of the spectral hardness.  The choice of peak photon flux (as opposed to energy fluence, energy flux, or total photon fluence) is based not on the physics of bursts (i.e., theories of fundamental burst properties) but on detector triggers:  most detectors trigger on a statistically significant increase in the count rate over a specified energy band.  The detection threshold is the minimum peak count rate for which the detector would have triggered.  Even though detectors trigger on the count rate in different energy bands, the peak count rate can be translated into the peak photon flux over a fiducial energy band with knowledge of the burst spectrum and the detector's energy response.  Note that the count rate is a detector-dependent quantity while the photon flux is an intrinsic description of the burst photons arriving in the Solar System. Therefore, the peak photon flux over a common energy band provides a convenient instrument-independent measure of burst intensity.  Because few bursts have lightcurves where the maximum flux is constant over seconds, the peak photon flux will depend on the accumulation time $\\Delta t$, the time resolution with which the flux is measured. Although most detectors trigger on a variety of accumulation times, $\\Delta t=1$~s is usually included in a detector's set of accumulation times.  In addition, a detector's sensitivity for one value of $\\Delta t$ can be translated into the average sensitivity for other values using an ensemble of burst lightcurves (Band 2002).  Therefore I use $\\Delta t=1$~s for this work.  Many studies have presented results using peak photon fluxes in the 50--300~keV band, principally because this was main trigger band of the Burst and Transient Source Experiment (BATSE) on the {\\it Compton Gamma-Ray Observatory (CGRO)}.  While any fiducial energy band can be used, particularly soft transient events, such as the recently-discovered X-Ray Flashes (XRFs, which may or may not be related to classical gamma-ray bursts---Heise et al. 2001), will produce little flux in the 50--300~keV band. Therefore I choose to use the 1--1000~keV band because most burst detectors operate within this broad band, and most bursts have $E_p$ in this band.  This study focuses predominantly on a comparison of the sensitivity of a number of detectors.  In \\S 2 I discuss the methodology used in this comparison, while \\S 3 presents the relevant information about each detector.  The results and their implications are discussed in \\S 4. ",
        "conclusions": "\\subsection{Detector Comparison} The maximum sensitivities for the different detectors are presented in Figures~2--9; the fraction of a detector's FOV at or near the maximum sensitivity varies from detector to detector. The sensitivity decreases away from the detector's normal because of the projection of the detector plane to the burst (proportional to the cosine of the inclination angle).  In the partially coded region of a coded mask detector the detector's walls shadow part of the detector plane.  In addition, the background induced by the particle flux, which is modelled crudely in my study, varies over an orbit, raising and lowering the sensitivity. Thus the curves in Figures~2--9 will shift up and down (mostly up) for different angles to the detector and for different parts of the spacecraft's orbit.  Nonetheless, the curves are a measure of the relative sensitivities of the different detectors.  The GBM's NaI detectors will be less sensitive than the BATSE LADs for bursts with $E_p>100$~keV because they will have much less area, although the GBM's low energy efficiency will give it a comparable sensitivity for bursts with low $E_p$. With significantly less area that the BATSE LADs, the FREGATE is less sensitive for bursts with $E_p>70$~keV. However, the FREGATE's low energy sensitivity (down to 6~keV) increases its sensitivity to detect particularly soft bursts relative to BATSE. Also, FREGATE triggers on a variety of $\\Delta E$ while BATSE triggered on only one value of $\\Delta E$. {\\it EXIST} will be more sensitive than Swift because the aperture flux (per detector area) is almost a factor of 2 smaller while the detector area is larger. Note that although the total detector area of just one of {\\it EXIST}'s three telescopes will be 27,000~cm$^2$, no point in the FOV is in the fully coded region of all 9 modules. These relative sensitivities are also affected by the trigger significance required of a rate increase.  The energy dependence of the detector efficiency affects the energy dependence of the sensitivity.  CZT has high efficiency below $\\sim 100$~keV and then decreases, while NaI's efficiency peaks at $\\sim 100$~keV and remains high until $\\sim1$~MeV.  Thus the sensitivity of the BATSE LADs (which were NaI detectors) and the GBM NaI detectors decreases significantly for $E_p<100$~keV, while the sensitivity of the CZT detectors (Swift and {\\it EXIST}) decrease much less for low $E_p$.  {\\it EXIST} will use thicker CZT detectors than Swift, which will increase the high energy efficiency, and thus {\\it EXIST}'s high energy sensitivity is greater than Swift's by more than would be predicted by the increase in area.  A lower low energy cutoff also increases the low energy sensitivity since more of the spectrum can be detected. This explains the comparable low energy sensitivities of FREGATE and Swift, even though Swift's effective area will be much larger than FREGATE's: FREGATE is sensitive above 6~keV and Swift's spectrum will begin at $\\sim15$~keV.  As expected, the trigger energy bands also affect the energy sensitivity.  The cusps evident in the curves for FREGATE, Swift, GBM and {\\it EXIST} result from high ($\\sim$50--150~keV) and low ($\\sim$15--50~keV) trigger bands.  Figure~10 shows the sensitivity of the four trigger energy bands proposed for Swift. \\subsection{Implications} Gamma-ray bursts will populate the $E_p$--$F_T$ plane shown in Figures 2--9.  Thus the detector sensitivities shown on these figures show which burst populations the detectors will detect. FREGATE is more sensitive below $E_p=100$~keV than BATSE, particularly for bursts without a high energy tail (i.e., for $\\beta<-3$). Thus BATSE did not trigger on X-Ray Flashes (XRFs), transients with low $E_p$ (many of the XRFs {\\it Beppo-SAX}'s WFC detected are untriggered events in the BATSE data---Kippen et al. 2002), whereas the FREGATE has detected XRFs. Mallozzi et al. (1995) showed that on average bursts become softer as they become fainter.  Kippen et al. show that the XRFs detected by {\\it Beppo-SAX} appear to be the low intensity extension of the Mallozzi et al. trend.  A comparison of the BATSE and Swift sensitivities (Figures~2 and~7) indicates that on the 1~s timescale Swift may not detect fainter bursts with $E_p>100$~keV, but it will detect soft bursts (e.g., XRFs) that are a factor of $\\sim 2$ fainter than BATSE's threshold.  Swift will trigger on timescales both shorter and longer than the 64~ms, 256~ms and 1.024~s timescales on which BATSE triggered, and thus will detect fainter bursts with certain types of lightcurves (e.g., short bursts)."
    },
    "0212/astro-ph0212178_arXiv.txt": {
        "abstract": "We have compared the intensity distribution of molecular line emission with that of dust continuum emission, and modeled molecular line profiles in three different preprotostellar cores in order to test how dynamical evolution is related to chemical evolution, and whether we can use different chemical tracers to identify specific dynamical evolutionary stages. We used dust continuum emission to obtain the input density and temperature structures by calculating radiative transfer of dust emission. Our results show that chemical evolution is dependent on dynamical processes, which can give different evolutionary timescales, as well as the density structure of the core. ",
        "introduction": "Low mass star formation has received more detailed study, both in theory and observation, compared to massive star formation. An evolutionary seqence of low mass star formation based on the spectral energy distribution (SED) has been developed (Adams et al. 1987). However, the initial conditions of collapse are still poorly understood. Preprotostellar cores (PPCs) are believed to be gravitational bound, but they have no central hydrostatic protostar (Ward-Thompson et al. 1994). Therefore, PPCs can be considered as the potential sites of future star formation and give us chances to probe the initial conditions of star formation. Evans et al. (2001) modeled dust continuum in three PPCs using Bonnor-Ebert spheres and showed that the dust temperature decreases toward the center because PPCs are heated only by the interstellar radiation field, and the radiation from the cores is optically thin. Therefore, they argued that the actual density structure of a preprotostellar core is more centrally condensed compared to the intensity distribution of the dust continuum emission.  Several studies on chemical evolution in PPCs have been done by Bergin \\& Langer (1997), Aikawa et al. (2001), Caselli et al. (2002), and Li et al. (2002). They tested different dynamical scenarios, but all included the interaction of gas phase molecules with the surfaces of dust grains as well as gas phase chemistry and showed that some molecules such as CS and CO are depleted substantially from the gas phase as the density increases. Those molecules have high binding energy onto the surfaces of dust grains, so their timescales for evaporation by thermal energy or cosmic rays are much longer than those of molecules that have lower binding energy. For instance, nitrogen bearing molecules such as N$_2$H$^+$ and NH$_3$, which become abundant in later stages of dynamical evolution and have low binding energy, do not show as significant depletion as CO and CS. Comparison of the abundance change of a molecule with time and the differential distributions of molecules in a given time seems to be able to trace the dynamical evolutionary stages. We tested this idea by comparing the line emission from several molecules that have different chemical evolution timescales with dust emission, which can trace physical conditions, in three PPCs: L1512, L1544, and L1689B.  We observed the three PPCs in C$^{18}$O ($\\rm J=2-1$ and $\\rm J=3-2$), C$^{17}$O ($\\rm J=2-1$), DCO$^+$ ($\\rm J=3-2$), $\\rm HCO^+$ ($\\rm J=3-2$), and $\\rm H^{13}CO^{+}$ ($\\rm J=3-2$) with the 10.4 m telescope of the Caltech Submillimeter Observatory (CSO) at Mauna Kea, Hawaii from 1995 to 2002. We also observed these cores in $\\rm H^{13}CO^{+}$ ($\\rm J=1-0$), $\\rm N_2H^+$ ($\\rm J=1-0$), and CCS ($\\rm N_J=4_3-3_2$) with the 45 m telescope of the Nobeyama Radio Observatory in Japan on January 2002. See Lee et al. (2003) for details about observations. ",
        "conclusions": "Our simple analysis and detailed modeling of molecular line profiles combined with dust continuum emission show that L1512 is least centrally condensed, but CCS, CO, and HCO$^+$ are significantly depleted. Even N$_2$H$^+$ is possibly depleted in L1512. On the other hand, L1689B has a similar density structure to L1544, which is quite centrally condensed and has substantial depletion of CCS, CO, and HCO$^+$, but no molecules except for CCS show significant depletion. This result may indicate that L1689B has evolved too fast to have time to evolve chemically. Therefore, we suggest that the stage of dynamical evolution of a core cannot be simply probed by its chemical status because chemical evolution depends on the size of dust grains (Caselli et al. 2002, in preparation) or the absolute dynamical timescale, which is different for each dynamical process, as well as its density structure that shows the relative dynamical evolutionary stages (Table 2).  \\begin{table} \\begin{center} \\caption{The results of this study}\\vspace{-1.0em} \\renewcommand{\\arraystretch}{1.0} \\begin{tabular}[h]{lccc} \\hline & L1512 & L1544 & L1689B \\\\ \\hline Density structure &  Young & Evolved & Evolved\\\\ Chemical evolution  &  Evolved & Evolved & Young\\\\ Timescale   &  $\\tau_{che}<<\\tau_{dyn}$ &$\\tau_{che}\\approx\\tau_{dyn}$ & $\\tau_{che}>>\\tau_{dyn}$\\\\ Dynamical state &Stable (?) &AD (?) & Free-fall (?)\\\\ \\hline \\\\ \\end{tabular} \\label{tab-single} \\end{center} \\end{table}"
    },
    "0212/astro-ph0212387_arXiv.txt": {
        "abstract": "{Many radio sources like quasars, blazars, radio galaxies, and micro-quasars exhibit circular polarisation (CP) with surprising temporal persistent handedness. As a possible explanation we propose that the CP is due to Faraday conversion (FC) of linear polarisation (LP) synchrotron light which propagates along a line-of-sight (LOS) through twisted magnetic fields. The rotational nature of accretion flows onto black holes naturally generates the required magnetic twist in the emission region, independent of whether it is a jet or an advection dominated accretion flow (ADAF). The expected twist in both types of flows is of the order of what is required for optimal CP generation. This scenario requires that Faraday rotation (FR) is insignificant in the emission region. Although this is an assumption, it relaxes constraints on the plasma parameters, that were given in scenarios which rely on FR, since there the strength of FR can not be too far from the optimum.  The proposed mechanism works in electron-positron ($\\e^\\pm$) as well as electron-proton ($\\e/{\\rm p}$) plasma. In the latter case, the emission region should consist of individual flux tubes with independent polarities in order to suppress too strong FR. The predominant CP is expected to mostly counter-rotate (rotation is measured here in sky-projection) with respect to the central engine in all cases (jet or ADAF, $\\e^\\pm$ or $\\e/{\\rm p}$ plasma).  If the proposed mechanism is indeed operating, it will allow to measure the sense of rotation of quasar engines.  The engine of SgrA$^*$ is then expected to rotate clockwise and therefore counter-Galactic, as do the young hot stars in its vicinity, which are thought to feed SgrA$^*$ by their winds.  Similarly, we expect the microquasars SS~443 and GRS~1915+105 to rotate clockwise. Generally, sources with Stokes-$V<0$ ($V>0$) are expected to rotate clockwise (counter-clockwise) in this scenario. ",
        "introduction": "}  The long lasting interest in CP from quasar-like systems \\citep[e.g.][ in the following J\\&O'D]{1977ApJ...214..522J} increased recently due to the detection of several new CP sources as SgrA$^*$ \\citep{1999ApJ...523L..29B, 1999ApJ...526L..85S}, M81$^*$ \\citep{2001ApJ...560L.123B}, the micro-quasars SS~433 and GRS~1915+105 \\citep{2000ApJ...530L..29F, 2002MNRAS.336...39F}, low luminosity AGNs \\citep{2002ApJ...578L.103B}, and many additional powerful quasars \\citep{1998Natur.395..457W, 1999AJ....118.1942H} and blazars \\citep{2001ApJ...556..113H}.  In the following we will summarise these systems under the general term quasar, assuming that a similar mechanism produces CP in most of them. The level of CP is $\\sim 1\\%$ and below, usually (but not in the case of SgrA$^*$) much lower than the level of LP.  On the one hand CP is highly time variable, but on the other hand in many sources it exhibits a very persistent rotational sense (per source and at nearly all frequencies where detected), which is constant on timescales of decades \\citep{1984MNRAS.208..409K, 1999AJ....118.1942H, 2002ApJ...571..843B}, although exceptions exist. This is far in excess of the dynamical timescale of the central quasar engine from which the emission originates.  In these sources synchrotron radiation from relativistic electrons produces LP, but only a very small and usually negligible amount of CP. A number of mechanisms have been proposed to explain the observed CP, among which FC of LP to CP seems to be the most likely process \\citep{2002MNRAS.336...39F,2002A&A...396..615M}. \\cite{2002A&A...388.1106B} and \\cite{2002ApJ...573..485R} (B\\&F and R\\&B in the following) provide a good introductions into the matter, a detailed discussion of the various CP generation mechanisms, and further references. For a discussion of scintillation models for CP generation \\cite{2002PASA...19...43M} should be consulted.  FC of synchrotron emission is a two-step process, since the emitted polarisation state has to be rotated before it can be Faraday converted into circular polarisation. This can happen by FR, as discussed by e.g. B\\&F and R\\&B and used as a starting point of this work, or it can be done by a systematic geometrical rotation of the magnetic field along the LOS, as originally proposed by \\cite{1982ApJ...263..595H} and discussed here.  The motivation for the standpoint adopted in this Letter, the assumption that the rotation of LP is a geometric effect rather than FR, is the observational fact that CP changes are very rare, much more rare than LP rotator events, and that a predominant CP sign seems to be persistent in many sources.  As explained in the following, this observational fact requires in the FR based models a constant magnetic polarity in the emission region over timescales far in excess of the dynamical timescale of the central engine, since the sign of CP depends on the magnetic field polarity in such models.  In the geometrical model discussed here, the CP sign is fully determined by the rotational sense of the central engine. Therefore the persistence of the predominant CP sign would -- in this picture -- be a natural consequence of angular momentum conservation. ",
        "conclusions": "Twisted magnetic flux can produce circularly polarised synchrotron emission due to FC. If FR is sufficiently suppressed the expected sense of rotation of the emission is expected to be opposite to the one of the central engine, independent of the question whether the emission results form a synchrotron jet, or from a synchrotron emitting ADAF, and independent of the polarity of the magnetic fields. The FR suppression can be -- as discussed by B\\&F and R\\&B -- due to a charge symmetric $\\e^\\pm$ plasma, in which FR is absent, or due to small-scale magnetic field polarity changes which produce cancelling FR contributions, but do not affect FC if the different flux tubes show alignment.  Such polarity changes could be present if the magnetic field is organised in many helically aligned flux tubes of independent magnetic field polarity which originated in differently magnetised patches of the accretion disk.   This model is not in conflict with other proposed generation mechanisms for CP (as e.g.  stochastic Faraday rotation with a small mean magnetic field as proposed by B\\&F and R\\&B), rather, it supplements them. Several of its strengths listed below are also present in other models. However, a few of the strengths (No. 1., 2., and 6.) seem to be slight advantages, potentially suggesting that it could be the dominant mechanism in several cases: \\begin{enumerate} \\item The sign of the CP does not depend on the presence of a mean field, but on geometrical properties of the flow pattern in the central engine of a quasar, which are fixed by the system's angular momentum. This may explain naturally the observed sign persistence of the CP over periods which exceed the dynamical timescales of the system by orders of magnitude. \\item The requirement that the rotation of the angle between LP and projected magnetic field is of the order one (not much higher, since over-rotation reduces CP, not much lower in order to have FC operating) within the optically thin part of the emission region, is naturally fulfilled by the geometrical properties of jet and ADAF flows. This requirement implies some level of fine-tuning in other models where the rotation is due to FR. \\item CP is still converted LP in this scenario, so that the variability of CP should exceed the one of LP as observed. \\item The mechanism can work in jets and in ADAFs. \\item It is able to produce CP in an $\\e^\\pm$ and in an e/p plasma. \\item The relatively large robustness of the mechanism to variations in the underlying source properties -- as composition of the Faraday converting plasma and nature of the emission region (jet or ADAF) -- may explain why CP appears in a large variety of very different systems, such as micro-quasars, low luminosity AGNs, and powerful quasars. \\item In cases where CP exceeds LP (e.g. SgrA$^*$), an outer Faraday screen, e.g. located in a mixing layer around an $\\e^\\pm$ jet, could provide the depolarisation \\citep[e.g.][]{2001ApJ...556..113H}. Alternatively, a $\\pi/2$ rotation of the sky-projected magnetic field throughout the visible emission region (the full optical thin jet, or the $\\tau \\le 1$ region) could remove any observable LP, but leave a measurable amount of CP. \\end{enumerate} There are conditions, under which the here proposed relation between the (predominant) rotational sense of CP and the rotation of the central engine is not valid anymore. Whenever the average FR is stronger than the twist in the magnetic fields along the LOS, the sign of the average magnetic field will determine the sign of CP as in the models of B\\&F and R\\&B. This may occur temporarily, if thermal plasma gets mixed into an $\\e^\\pm$ jet. If the emission region contains an e/p plasma, the occurrence of a mean magnetic field -- either as a statistical fluctuation, or due to some physical reason -- can increase the average Faraday depth and lead to a CP reversal. External CP generation from LP due to environmental scintillation also can mask the intrinsic CP signature of a quasar. Also within the scenario in which CP is due to bending in magnetic fields, CP sign reversals are possible, whenever the magnetic twist within the emission region is not linked to the central engine rotation. This may happen temporarily in violent shocks in a jet outflow, and would suggest that CP sign reversals of this origin are likely accompanying emission outbursts of the source\\footnote{This may explain the observed CP sign reversal in GRO J1655-40 in during an outburst in 1994, as reported by \\cite{2002A&A...396..615M}}.  The CP should also change if the rotational sense of the central engine is changed, e.g. due to a black hole merging event, or due to a massive infall of fresh material onto the accretion disc.  Note, that the rarity of CP sign reversals indicate that the above discussed conditions should be exceptional, if CP is due to the here proposed mechanism.   The proposed mechanism provides testable predictions: \\begin{enumerate} \\item The sign of CP from a quasar measures the rotation direction of the quasar's engine. The electric vector of radio emission should rotate retrograde in the sky-plane with respect to the engine rotation. Positive $V$ polarisations implies therefore positive $=$ counter-clockwise rotation (on the sky) of the engine. \\item Thus, we expect SgrA$^*$, which exhibited $V<0$ during the last 20 years, to rotate clockwise. This is retrograde with respect to the galactic rotation and the rotation of the molecular gas cloud in the galactic centre. However, it has the same rotation sense as the population of young hot HeI stars in its direct vicinity \\citep{2000MNRAS.317..348G}, which were proposed to feed SgrA$^*$ via their stellar winds \\citep[][ for discussion and references]{2001dscm.conf..291G}. Conservation of angular momentum in this wind should lead to a retrograde (clockwise) rotating accretion disk. Alternatively, the spin of the central black hole may determine the twist of the magnetic fields in the emission region. \\item We expect the microquasars GRS~1915+105 and SS~433 (both exhibit $V<0$) also to rotate clockwise, and we could give corresponding expectations for the other sources with detected CP. \\item The predominant sign of CP should be temporally constant. If the model is correct, quasar-like CP sources should only exhibit CP flips under exceptional circumstances, as discussed above. \\item CP from the counter jet should have opposite handedness than CP from the approaching jet. This prediction may allow a discrimination of the presented model from FR based scenarios, since in the latter the same predominant CP sign is expected from both jets for a dipole-like mean field component. \\item There are possibly also spectral signatures, which allow to discriminate the different models. E.g. in FR based scenarios frequency dependent LP oscillations can occur, as illustrated by model 2 in Fig. \\ref{fig:example}, whereas this is impossible in the magnetic twist based scenarios. However, due to cancellation effects in the superposition of many slightly different LOSs such oscillations likely appear as a strong suppression of LP. Due to the high complexity of the resulting polarisation spectra general rules are not obvious and would require further investigations. In that context it is interesting to note that Beckert (in preparation) recently presented new model calculations for SgrA$^*$. In order to simultaneously explain new high frequency LP measurements of SgrA$^*$ model parameters seem to be favoured for which magnetic twist is more important for CP production than FR. \\end{enumerate} If the scenario of CP production proposed here could be demonstrated to be operatively in quasars, it would give us a tool to measure the sense of rotation of the most powerful and enigmatic engines of the universe -- the violent matter flows in the direct vicinity of black holes -- by simply looking at the spin of the received photons."
    },
    "0212/astro-ph0212309_arXiv.txt": {
        "abstract": "Both preheating of the intergalactic medium and radiative cooling of low entropy gas have been proposed to explain the deviation from self-similarity in the cluster $L_{\\rm X}-T_{\\rm X}$ relation and the observed entropy floor in these systems. However, severe overcooling of gas in groups is necessary for radiative cooling alone to explain the observations. Non-gravitational entropy injection must therefore still be important in these systems. We point out that on scales of groups and below, gas heated to the required entropy floor cannot cool in a Hubble time, regardless of its subsequent adiabatic compression. Preheating therefore shuts off the gas supply to galaxies, and should be an important global feedback mechanism for galaxy formation. Constraints on global gas cooling can be placed from the joint evolution of the comoving star formation rate and neutral gas density. Preheating at high redshift can be ruled out; however the data does not rule out passive gas consumption without inflow since $z \\sim 2$. Since for preheated gas $t_{\\rm cool} > t_{\\rm dyn}$, we speculate that preheating could play a role in determining the Hubble sequence: at a given mass scale, high $\\sigma$ peaks in the density field collapse early to form ellipticals, while low $\\sigma$ peaks collapse late and quiescently accrete preheated gas to form spirals. The entropy produced by large scale shock-heating of the intergalatic medium is significant only at late times, $z<1$, and cannot produce these effects. ",
        "introduction": "Numerical simulations of dark matter halos reveal an almost universal density profile over a wide range of masses \\cite{NFW,ghigna00,klyp01,power02}. If the gas faithfully traced the dark matter, we would expect the bolometric X-ray luminosity to scale with the cluster virial temperature as $L_{\\rm X} \\propto T^{2}$ \\cite{kaiser86}, while observations indicate significant steepening of this relation for low-temperature clusters \\cite[and references therein]{pon00}, suggesting that some non-gravitational process is affecting the properties of the gas. Indeed, observations of the X-ray surface brightness profiles in poor groups reveal the existence of an entropy floor, $S=T/n_{\\rm e}^{2/3} \\sim 100 {\\rm keV cm^2}$ \\cite{ponmanetal}, in the gas. Such an entropy floor reduces the central gas density and breaks the self-similarity of gas density profiles. There are two schools of thought as to how this entropy floor could have arisen: (i) as a result of non-gravitational heating of order $\\sim 1$keV per particle, arising from winds driven either by supernovae or active galactic nuclei (e.g., \\pcite{kaiser,ponmanetal}), (ii) as a result of radiative cooling of the low entropy gas, leaving behind only the high entropy gas (which cannot cool within a Hubble time; e.g., \\pcite{voitbryan}). Numerical simulations which incorporate either of these schemes can reproduce the observed cluster $L_{\\rm X}-T$ relation.  In this paper, we critique the first scenario, on the basis of the disastrous effects the required level of preheating would have on subsequent galaxy formation. Groups and clusters are the most massive non-linear structures present today; to have a significant effect on their gas profile, a large amount of energy must be injected ($\\sim 1$keV per particle). The virial temperature of galaxies is much smaller (for an ${\\rm L_{*}}$ galaxy, typically $T_{\\rm vir} \\sim 0.1$keV); galaxies will therefore be unable to accrete gas heated to such high entropies. Furthermore, we shall show that the gas they do manage to accrete cannot cool within a Hubble time. The net effect would be a sharp downturn in the comoving star formation rate. \\scite{momao} have also recently examined the impact of preheating on galaxy formation; they find that preheating could have a strong impact on the X-ray luminosity and morphology of galaxies. However (unlike the present work), they do not consider if preheating will suppress subsequent star formation. Throughout, we assume a cosmology where $(\\Omega_{\\rm m},\\Omega_{\\Lambda},\\Omega_{\\rm b} h^{2},h,\\sigma_{8 h^{-1}})=(0.3,0.7,0.019,0.7,0.9)$, consistent with current observational constraints (e.g. \\pcite{cmb,hstkey,smith02}). ",
        "conclusions": "\\subsection{Shock heating of the Warm Hot Intergalactic Medium} Cosmological simulations predict that $\\sim 30-40 \\%$ of present day baryons reside in a warm-hot intergalactic medium (WHIM), with temperatures ${\\rm 10^{5} < T < 10^{7} K}$, and mean overdensities $\\delta \\sim 10-30$, rather than in virialized objects \\cite{CO99,Croftetal,Daveetal}. Can gas heated in such large scale shocks accrete and cool in galaxies?  \\scite{Daveetal} find that the WHIM at z=0 approximately follows the equation of state: $\\rho/\\rho_{\\rm b}= T/10^{4.7}{\\rm K}$. We can use this to estimate the entropy of the WHIM: \\begin{equation} S_{\\rm WHIM}(z=0) \\approx 340 \\left( \\frac{\\delta}{30} \\right)^{1/3} {\\rm keV \\, cm^{-2}}. \\end{equation} Since $S_{\\rm WHIM} > S_{\\rm crit}$, none of the gas in the WHIM, even if accreted onto galaxies, can cool in a Hubble time. Thus, shock heating by large scale structure could conceivably play a similar role as preheating in raising the entropy of the IGM and suppressing accretion and gas cooling.  However, in hierarchical structure formation scenarios, such shock heating cannot be responsible for the observed entropy floor in groups and clusters. Shock heating of the WHIM corresponds to $\\sim 1 \\sigma$ fluctuations turning nonlinear today; the cores of groups and clusters were assembled at high redshift when they were $ > 2 \\sigma$ fluctuations. Gas shocked in the WHIM phase is likely shocked again as it is boosted to the higher entropies $S \\sim 1350 (T/1 {\\rm keV}) (n/2 \\times 10^{-5}{\\rm cm^{-3}})^{-2/3} {\\rm keV \\, cm^{-2}}$ associated with the outskirts of present-day groups and clusters. Shock heating of the WHIM is likely only to affect gas accretion in voids, where the halos have shallow potential wells and are unable to accrete gas of high entropy.  We can estimate the redshift evolution of the WHIM entropy as follows. The postshock temperature of the WHIM can be estimated as: \\begin{equation} T_{\\rm WHIM}(z) = K (H L_{\\rm nl})^{2}, \\end{equation} where $H(z)$ is the Hubble constant, $L_{\\rm nl}$ is the nonlinear length scale in proper coordinates; this produces a good fit to the results of numerical simulations \\cite{CO99,Daveetal}. We have chosen the normalization constant $K$ so as to reproduce the median temperature for a given overdensity obtained by \\scite{Daveetal} at z=0. At a given overdensity $\\delta$, the entropy of the WHIM scales as $S_{\\rm WHIM} \\propto H(z)^{2} L_{\\rm nl}^{2} (1+z)^{-2}$; it thus falls rapidly with redshift. We plot $S_{\\rm WHIM}(z)$ in Figure (\\ref{Fig:S_WHIM}); it only exceeds $S_{\\rm crit}$ at $z < 1$. In practice there will be significant scatter about the median entropy of the WHIM (for instance, the temperature distribution broadens at high redshift; see Fig. 4 of \\pcite{Daveetal}). It would be interesting to use existing numerical simulations to compute the distribution function of entropy in the WHIM, as well as the redshift evolution of the median IGM entropy; such quantities to date have not been calculated.  \\begin{figure} \\psfig{file=WHIM_entropy.ps,width=80mm} \\caption{The evolution of the median entropy of the WHIM as a function of redshift. Note that $Sq_{\\rm WHIM} \\propto \\delta^{1/3}$; here we assume $\\delta\\sim 30$. Although in practice there will be wide scatter about this relation, we nonetheless see that the entropy of the WHIM falls very rapidly with redshift. In particular, it only exceeds the critical entropy $S_{\\rm crit}$ required to prevent cooling within a Hubble time at late times.} \\label{Fig:S_WHIM} \\end{figure}  It is amusing to speculate on the effect of a cosmological constant on the future star formation history of our universe. The growth function is given by \\cite{Heath}: \\begin{equation} D_{1}(a) \\propto H(a) \\int_{0}^{a} \\frac{\\dd a^{\\prime}}{a^{\\prime 3} H(a^{\\prime})^{3}}, \\end{equation} where $a=1/(1+z)$ is the scale factor and $H(a)=(\\Omega_{\\rm m}a^{-3} + \\Omega_{\\rm k}a^{-2} + \\Omega_{\\Lambda})^{1/2}$, where $\\Omega_{\\rm k} = 1- \\Omega_{\\rm m} - \\Omega_{\\Lambda}$. Note that at early times when $a$ is small, $D_{1}(a) \\propto a$. If we normalize this so that $a=1$ and $D_{1}(0)=1$ today, then in the future when $a \\rightarrow \\infty$, $D_{1}(a) \\rightarrow 1.39$. Thus, structure formation is already freezing out; the non-linear mass $M_{*}$ will increase by at most a factor of 4 from its present day value, in contrast to its rapid growth in the past. Thus, gravitational shock heating is slowing down and becoming unimportant, as the universe enters a period of exponential expansion. Gas in already virialized halos will eventually be able to cool as the universe ages, modulo feedback by star formation and AGN activity fueled by gas cooling. This is in contrast to the $\\Omega_{\\rm m}=1$ SCDM case, when $D_{1}=a$ and structure formation and gravitational shocks continue apace; in this case most of the non-stellar baryons in the universe would end up in a hot phase which will never cool.  \\subsection{Could Preheating determine Galaxy Morphology?}  We speculate that preheating may be responsible for the origin of the Hubble sequence. Gas in dark matter halos which form before the epoch of preheating are assumed to form an elliptical galaxy in a rapid collapse, since halos below group scales have gas with cooling times shorter than their dynamical time $t_{\\rm cool} < t_{\\rm ff}$ \\cite{reesostriker}. On the other hand, halos which collapse after preheating will accrete gas with entropies more characteristic of groups and clusters; in particular, they can cool their gas only slowly, with $t_{\\rm cool} > t_{\\rm ff}$. Cooling gas in such halos may instead form a disk galaxy. In fact, \\scite{momao} point out that preheating could potentially solve two problems associated with disk galaxies: explaining the anomolously low X-ray surface brightnesses associated with their halos (since preheating decreases gas densities and X-ray emissivities), and preventing angular momentum transport from the disk to the dark matter during disk formation (because the gas distribution is more extended than that of the dark matter, it retains a higher specific angular momentum during mergers). We note in passing that the low X-ray luminosities of halos associated with disks is perhaps unsurprising: a simple extrapolation of the $L_{\\rm X}-T_{\\rm X}$ relation to halo temperatures typical for the hosts of disks ($\\sim 0.1$keV) gives X-ray luminosities consistent with observations (few $\\times 10^{40} \\, {\\rm erg \\, s^{-1}}$). Thus, whatever physics is suppressing the X-ray luminosities of group operates in disk galaxies too.  In a hierarchical formation scenario, the progenitors of massive clusters typically formed earlier than those of lower mass systems. Thus, massive clusters contain a larger fraction of galaxies which formed prior to the epoch of preheating; the fraction of elliptical galaxies should increase as the mass of the dark matter halo increases. To compute this fraction of galaxies which formed before preheating (and therefore formed ellipticals) we create merger trees for halos of different masses at $z=0$. We then identify when $10^{12}M_\\odot$ progenitor halos form in these trees. If the progenitor formed before preheating (assumed to occur at $z=2$ here) the halo is flagged as hosting an elliptical galaxy, while halos which collapse later are assumed to host a disk galaxy. In Figure~\\ref{fig:morph} we show the fraction of galaxies formed in $10^{12}M_\\odot$ halos which are ellipticals as a function of the mass of the halo in which they are found at $z=0$. This is compared against the observational determinations of \\scite{balogh02} of the elliptical fraction (shown as circles), with their estimate of the ``field'' elliptical fraction shown as a dashed line. Our simple calculation fits rather well with these data (note that the ``field'' should correspond approximately to $M_*\\sim 10^{13}h^{-1}M_\\odot$ halos). Since the typical mass of halos increases in denser environments we expect this model to produce the observed morphology-density relation, at least qualitatively. A direct comparison requires a detailed model of galaxy formation, since it requires calculating the local galaxy density within the model.  \\begin{figure} \\psfig{file=morphology.ps,width=80mm} \\caption{The fraction of galaxies formed in $10^{12}M_\\odot$ halos which are ellipticals, as a function of the mass of the halo in which they are found at $z=0$. Points with errorbars show the determinations of \\protect\\scite{balogh02}. (The X-ray luminosities of clusters in \\protect\\scite{balogh02} were converted to cluster virial masses using the relation of \\protect\\scite{rei99}.) The dashed line shows the ``field'' elliptical fraction from \\protect\\scite{balogh02}, with dotted lines showing their errorbars.} \\label{fig:morph} \\end{figure}  This is of course little more than a consistency check, rather than proof of principle. The idea that ellipticals collapsed at much higher redshift than spirals is an old one. It was suggested long ago that there may be a direct correspondence between peak height and galaxy morphology, with ellipticals forming from rare ($\\sim 3 \\sigma$) density fluctuations, while spirals form from more common ($\\sim 2 \\sigma$) ones \\cite{blumenthal,evrard}; such a scheme can reproduce their observed relative abundances fairly well and would also explain the increased clustering strength of ellipticals. Empirically, it is known that ellipticals exhibit remarkably tight correlations in various global properties and have very small inferred age dispersion which indicate a high formation redshift $z \\sim 2-3$, unless their formation was synchronized to an implausible degree. The evidence includes the tightness of the fundamental plane \\cite{vandok}, and the extreme homogeneity of their optical colors \\cite{boweretal,ellisetal}. However, the physical motivation for early formation of ellipticals rather than spirals has been rather weak.  The interesting feature that preheating introduces is a physical reason why there should be a division in morphology between early and late collapse of halos of the same mass, because of the change in the cooling properties of the gas. There may be some critical level of entropy (likely $t_{\\rm cool}(S_{\\rm crit}) \\sim t_{\\rm dyn}$) where galaxy formation switches from one mode to another. To test these speculations, it would be necessary to conduct high resolution simulations of a single galaxy, to observe how galaxy morphology and disk structure changes as the entropy of the accreted gas is increased. To some extent, we know already from previous simulations (based on suggestions that photo-ionization could delay cooling until $z \\sim 1$) that feedback processes which delay gas cooling result in less angular momentum transfer and more disk-like features \\cite{weiletal}.  Preheating may also help alleviate possible problems with the dearth of galaxies found in voids \\cite{peebles,mathis,bensonvoids}. Since the distribution of halo masses in voids is shifted to lower masses relative to the field, preheating will be particularly effective in suppressing galaxy formation in these environments. However, lower mass halos also tend to collapse earlier, so a larger fraction of these halos will form prior to preheating. This latter effect seems to dominate. Using dark matter halos from the simulations of \\scite{bensonvoids}, we identified halos which formed after preheating at $z=2$ (assuming the distribution of formations redshifts given by \\scite{verde} and no correlation of formation time with environment). Removing these halos from the sample makes only minor differences to the void probability function for example. Although these simulations have a rather poor mass resolution for the study of void galaxies, we expect the effects of preheating to become even less for lower mass halos, since an even greater fraction of these should form prior to preheating.  \\subsection{Conclusions}  In this paper, we stress that radiative cooling alone cannot account for the observed entropy floor in clusters and the minimum entropy required to produce the observed deviation from self-similarity in the $L_{\\rm X}-T_{\\rm X}$ relation. This is only possible with severe overcooling. Thus, some form of entropy injection--perhaps through supernovae or AGN winds--seems necessary (though radiative cooling still plays a very important role in reducing the energetic requirements for producing an entropy floor \\cite{voitetal}). The level of entropy injection required to produce observable changes in the density profiles of deep potential wells of groups and clusters must have a drastic effect on smaller halos which host galaxies. We show that entropy injection can indeed play a very important role in regulating gas accretion and cooling in galaxies, since gas heated to the entropy floor can never cool in a Hubble time, regardless of the densities it is compressed to: for gas temperatures typical of galaxy halos the cooling time depends almost exclusively on entropy and is relatively independent of density. Thus, after the epoch of preheating, the gas supply to galaxies was cut off. This is not inconsistent with observations if an epoch of preheating at $z \\sim 2$ is assumed. We speculate that at a critical level of entropy $t_{\\rm cool}(S_{\\rm crit}) \\sim t_{\\rm ff}$, the mode of gas collapse changes from free-fall and fragmenation (which produces ellipticals) to quiescent accretion (which produces spirals).  There are many aspects of the impact of preheating on galaxy formation which would be worth pursuing in greater detail, particularly with numerical simulations. One would be the redshift and spatial dependence of entropy injection--we have only considered a uniform level of entropy injection which appears instantaneously at $z \\sim 2$. Naively, we would expect entropy injection to be proportional to the integrated star formation, which is proportional to the local stellar density and local metallicity. Similarly, the spatial variation and distribution function of entropy due to shock heating of the IGM deserves attention, and the possibility that the entropy of the IGM may be responsible for galaxy morphology. It may well be that the entropy of the IGM regulates the amount of cooling and star-formation at any given cosmological epoch, and is the chief mechanism which prevents overcooling at high redshift."
    },
    "0212/astro-ph0212415_arXiv.txt": {
        "abstract": "Radio observations were used to detect the `active' galaxy population within rich clusters of galaxies in a non-biased manner that is not plagued by dust extinction or the K-correction.  We present wide-field radio, optical (imaging and spectroscopy), and ROSAT All-Sky Survey (RASS) X-ray data for a sample of 30 very rich Abell ($R \\ge 2$) cluster with $z \\le 0.25$.  The VLA radio data samples the ultra-faint radio ($L_{1.4} \\ge 2\\times 10^{22}$ W Hz$^{-1}$) galaxy population within these extremely rich clusters for galaxies with $M_{\\mathrm{R}} \\le -21$. This is the largest sample of low luminosity 20\\,cm radio galaxies within rich Abell clusters collected to date.  The radio-selected galaxy sample represents the starburst (Star formation rate $\\ge$ 5\\,M$_{\\sun}$ yr$^{-1}$) and active galactic nuclei (AGN) populations contained within each cluster. Archival and newly acquired redshifts were used to verify cluster membership for most ($\\sim 95$$\\%$) of the optical identifications. Thus we can identify all the starbursting galaxies within these clusters, regardless of the level of dust obscuration that would affect these galaxies being identified from their optical signature.  Cluster sample selection, observations, and data reduction techniques for all wavelengths are discussed. ",
        "introduction": "Optical studies of rich clusters to determine the star formation rates (SFR) of their cluster members have been ongoing since the early photometric study by \\cite{but78a} over 20 years ago. Such studies have sought to recover the star formation properties of galaxies in rich clusters via optical photometry, spectroscopy and more recently from {\\it Hubble Space Telescope} (HST) morphological studies. The discovery by \\citeauthor{but78a} (\\citeyear{but78a}; \\citeyear{but84b}) of the increased fraction of blue galaxies within the core of selected rich regular clusters indicated significant evolution in the SFR and/or galaxy type within the central regions of these massive systems over the last $\\sim 5$ Gyr (assuming H$_0$ = 75 km s$^{-1}$, q$_0$ = 0.1) or $z \\lesssim 0.4$. The change in SFR of galaxies in rich clusters may be related to hierarchical large-scale-structure (LSS) formation and the accretion of field galaxies.  Where the starburst population occurs within the cluster environment yields clues to the origin of the perturbing mechanism. To study this `active' galaxy population (starburst/AGN) requires a method that will detect this population in a unbiased manner.  Radio observations allow one to detect starbursts and AGN without the drawbacks of optical observations, such as dust attenuation and the optical K-correction. In principle, far-IR (FIR) observations would detect re-emission from the dust, of the UV photons emitted by the OB stars formed in a massive starburst. Until the launch of SIRTF, FIR instruments will not have the sensitivity or the resolution to be effective. However, the well known FIR-radio continuum correlation based on ongoing star formation allows easier detection of starbursting systems in distant clusters. Studies of the local radio luminosity function \\citep[\\eg][]{con91}, indicate that the starbursting and spiral galaxy population dominate the 20\\,cm radio luminosity function below 10$^{23}$ W Hz$^{-1}$.  Radio selection of cluster members by their low radio luminosity yields two kinds of objects - weak AGN and objects with active star formation. The FIR-radio correlation ties the nonthermal radio emission to the ongoing star formation \\citep{con92}. In starbursts, radio emission fades faster than blue starlight, so clusters with large number of radio sources will preferentially be those that have had a large numbers of recent or ongoing starbursts. This may indicate the entry of a new group of galaxies into the cluster with star formation in its members.  This paper is the first in a series studying the ultra-faint radio-selected galaxy populations associated with rich clusters of galaxies as a function of redshift. This paper presents the radio, optical, and X-ray data for a sample of 30 very rich clusters (Abell richness class $R \\ge$2) with $z \\le 0.25$ selected from the catalog of \\citet*[][hereafter ACO]{abe89}.  The VLA offers a large field of view ($\\sim$30 arcmin) permitting us to go out to 2.5 Mpc radius from the cluster core. This allows us to detect infalling starbursting field galaxies and any supercluster induced effects such as cluster-cluster mergers. Our detection limit of 2$\\times$10$^{22}$ W Hz$^{-1}$ allows detection of galaxies undergoing moderate to vigorous star formation with star formation rates (SFR) $\\ge$ 5\\,M$_{\\sun}$ yr$^{-1}$.  In this data paper, we discuss selection of the cluster sample along with the multi-wavelength observations and the data reduction techniques for a sample of the richest Abell clusters from 0.02 $\\le$ z $\\le$ 0.25. These clusters were chosen to be as rich or richer than our four rich $z = 0.4$ clusters (Abell 370, Cl0024$+$1654, Cl0939$+$4713, and CL1447$+$2619), which will be presented in future papers. We also present a radio-selected sample of starburst and AGN galaxies within these 30 rich clusters. Most of the detected galaxies had followup spectroscopy to verify cluster membership.  Future papers analyze the radio properties \\citep[Paper II;][]{mor00}, the optical properties of the clusters \\citep[Paper III;][]{mor00a}, and spectroscopy of the optical counterparts \\citep[Paper IV;][]{mor00b}.  Deep (15 hr) wide-field VLA B-array observations of Cl0939+4713 at $z = 0.41$ and subsequent spectroscopic analysis of the radio-selected galaxy population will appear in Paper V \\citep{mor00c}. HST NICMOS and WFPC2 results of several of the radio-selected galaxies in Cl0939+4713 were discussed in \\cite{sma99}. ",
        "conclusions": "We present the largest sample of low luminosity ($L_{1.4 GHz} \\ge 2\\times 10^{22}$ W Hz$^{-1}$) radio galaxies within extremely rich Abell clusters obtained to date. The radio observations allow us to probe the `active' galaxy population of these clusters without the difficulties experienced in optical observations, such as dust obscuration and the optical K-correction. The radio-selected galaxy sample represents the starburst (SFR $\\ge$ 5\\,M$_{\\sun}$ yr$^{-1}$) and AGN populations contained within each cluster. New redshifts were used to verify cluster membership for most of the optical identifications.  A total of 165 radio-selected galaxies have been detected within 2.5 Mpc of the centers of rich Abell clusters with $z \\le 0.23$ and $M_{\\mathrm{R}} \\le -21$. There are 89 radio-selected galaxies with luminosities between $2\\times 10^{22}$ W Hz$^{-1} \\le L_{1.4} \\le 10^{23}$ W Hz$^{-1}$.  There are 76 radio-selected galaxies with $L_{1.4} \\ge 10^{23}$ W Hz$^{-1}$. We have obtained 114 new spectra for clusters with $z \\le 0.23$. A large fraction (95$\\%$) of the optical identifications have measured redshifts.  In addition, new optical and X-ray data are presented for each of these clusters.  We describe the selection of the cluster sample, observations and the data reduction techniques for all three wavelengths, including optical spectroscopy."
    },
    "0212/astro-ph0212286_arXiv.txt": {
        "abstract": "The currently most popular models for the dynamical evolution of star clusters predict that the power-law Cluster Luminosity Functions (CLFs) of young star cluster systems will be transformed rapidly into the universal Gaussian CLFs of old Milky Way-type ``globular'' cluster systems.  Here, we provide the first evidence for a turn-over in the intermediate-age, approximately 1 Gyr-old CLF in the center of the nearby starburst galaxy M82, which very closely matches the universal CLFs of old Milky Way-type globular cluster systems.  This provides an important test of both cluster disruption theories and hierarchical galaxy formation models.  It also lends strong support to the scenario that these young cluster systems may eventually evolve into old Milky Way-type globular cluster systems.  M82's proximity, its shortest known cluster disruption time-scale of any galaxy, and its well-defined peak of cluster formation make it an ideal candidate to probe the evolution of its star cluster system to fainter luminosities, and thus lower masses, than has been possible for any galaxy before. ",
        "introduction": "For old globular cluster systems, with ages in excess of $\\sim 10^{10}$ yr, the shape of the cluster luminosity function (CLF) is well-established: it is roughly Gaussian, with an apparently universal peak luminosity, $M_V^0 \\simeq -7.4$ mag, and a Gaussian FWHM $\\sim 3$ mag (e.g., Whitmore et al.  1995, Harris 1996, 2001, Ashman \\& Zepf 1998, Harris et al.  1998).  On the other hand, the well-studied young, $\\lesssim 2$ Gyr-old star cluster population in the Large Magellanic Cloud displays a power-law CLF of the form $N(L) {\\rm d} L \\propto L^{\\alpha} {\\rm d} L$, where $L$ is the cluster luminosity and $-2.0 \\lesssim \\alpha \\lesssim -1.5$ (e.g., Elson \\& Fall 1985, Elmegreen \\& Efremov 1997).  {\\sl Hubble Space Telescope (HST)} observations are continuing to provide an increasing number of luminosity distributions for additional samples of young and intermediate-age compact star clusters in more distant galaxies (e.g., Whitmore \\& Schweizer 1995, Schweizer et al. 1996, Miller et al.  1997, Zepf et al.  1999, de Grijs, O'Connell \\& Gallagher 2001, Whitmore et al.  2002).  Although a large number of studies have attempted to detect a turn-over in young or intermediate-age CLFs, the shapes of such young and intermediate-age CLFs have thus far been consistent with power laws down to the observational completeness thresholds.  Based on the observational evidence discussed above, the currently most popular globular cluster formation models suggest that the distribution of the initial cluster luminosities and, therefore, the corresponding distribution of their initial masses is closely approximated by a power law (Harris \\& Pudritz 1994, McLaughlin \\& Pudritz 1996, Elmegreen \\& Efremov 1997, Ashman \\& Zepf 2001) of the form d$N(M) {\\rm d}M \\propto M^{\\alpha} {\\rm d}M$, where $-2.0 \\lesssim \\alpha \\lesssim -1.5$. Several processes are proposed for the transformation of the CLF of young star cluster systems into the Gaussian or log-normal distributions characteristic of old globular cluster-type distributions.  These include the preferential depletion of low-mass clusters both by evaporation due to two-body relaxation and by tidal interactions with the gravitational field of their host galaxy, and the preferential disruption of high-mass clusters by dynamical friction.  However, the relative importance of these disruption processes is still controversial (see, e.g., Vesperini [2000, 2001], and Bromm \\& Clarke [2002] for models that do not implicitly assume a power-law initial mass distribution).  If the age range within a given cluster system is a significant fraction of the system's mean age, fading due to aging of the non-coeval cluster population also severely affects the shape of its CLF: the observed CLF will de dominated by the younger clusters, because a fraction of the older clusters in the system will have faded to below the observational detection limit, thus introducing further incompleteness effects.  The detection of a characteristic turn-over luminosity in young or intermediate-age CLFs would lend strong support to the popular (but thus far only speculative) scenario that the young star clusters observed in mergers are the analogues to the ubiquitous old globular clusters at younger ages. ",
        "conclusions": "Thus, here we have presented the first conclusive evidence for a clear turn-over in the CLF of a 1 Gyr-old, roughly coeval cluster population. The CLF shape and characteristic luminosity is nearly identical to that of the apparently universal CLFs of the old globular cluster systems in the Galaxy, M31, M87, and old elliptical galaxies (e.g., Whitmore et al. 1995, Harris 1996, 2001, Ashman \\& Zepf 1998, Harris et al.  1998). This is likely to remain virtually unchanged for a Hubble time.  We have also shown that with the very short characteristic cluster disruption time-scale governing M82 B, its cluster mass distribution will evolve towards a higher characteristic mass scale than for the Galactic globular clusters by the time it reaches a similar age.  We argue, therefore, that this evidence, combined with the similar cluster sizes (de Grijs et al.  2001), lends strong support to a scenario in which the current M82 B cluster population will eventually evolve into a significantly depleted old Milky Way-type globular cluster system dominated by a small number of high-mass clusters.  This implies that globular clusters, which were once thought to be the oldest building blocks of galaxies, are still forming today in galaxy interactions and mergers.  However, they will likely be more metal-rich than the present-day old globular cluster systems.  This connection between young or intermediate-age star cluster systems and old globular clusters lends support to the hierarchical galaxy formation scenario. Old globular clusters were once thought to have been formed at the time of, or before, galaxy formation, i.e., during the first galaxy mergers. However, here we have shown that the evolved CLF of the compact star clusters in M82 B most likely to survive for a Hubble time will probably resemble the high-mass wing of the ``universal'' old globular cluster systems in the local Universe. Proto-globular cluster formation thus appears to be continuing until the present.  In order to better constrain the future evolution of the M82 B star cluster system it is important to consider a range of models, characterized by fewer or perhaps different orbital restrictions (e.g., Baumgardt 1998, Vesperini 2000, 2001).  This is, however, beyond the scope of the current {\\it letter} and will be done in a subsequent paper (de Grijs et al., in prep.).  \\paragraph{Acknowledgements} - We acknowledge useful and interesting discussions with R.W.  O'Connell and S.M.  Fall.  We also thank A. Helmi for a critical reading of the manuscript, and the anonymous referee for helpful suggestions.  RdeG acknowledges support from the Particle Physics and Astronomy Research Council (PPARC) and from The British Council under the {\\sl UK--Netherlands Partnership Programme in Science}."
    },
    "0212/astro-ph0212553_arXiv.txt": {
        "abstract": "The arguments in favor of the unified formation mechanism for both slow (Lynden-Bell's) bars and common fast bars are given. This mechanism consists in a certain instability that is akin to the well-known radial orbit instability; it is caused by the mutual attraction and alignment of axes of precessing star orbits (up to now, such a way of formation was considered only for slow bars).  The general theory of the low-frequency modes of a disk consisting of precessing orbits (at different angular velocities) is presented. The problem of determining these modes is reduced to the integral equations of a rather simple structure. The characteristic pattern speeds ($\\Omega_p$) of the low-frequency modes are of order of the mean orbit precession speeds ($\\bar\\Omega_{pr}$). The bar-modes also belong to this type of modes. The slow bars have $\\Omega_p \\approx \\bar\\Omega_{pr}$; for the fast bars, $\\Omega_p$ may far exceed even the maximum precessing speed of disk orbits (however, $\\Omega_p$ remains to be of order of these precessing speeds). The possibility of such an excess of $\\Omega_p$ over $\\Omega_{pr}^{\\max}$ is connected with the effect of ``repelling'' orbits that tend to move in the direction opposite to that they are being pushed.  The preliminary analysis of the orbit precession patterns for a number of typical potentials is given. It is noted that the maximum radius of the ``attracting'' circular orbits ($r_c$) may be used as a reasonable estimate of a bar length. ",
        "introduction": "It is commonly supposed that a galactic bar can belong to one of two types --- either common fast bars or Lynden-Bell slow bars (see, e.g., Sellwood 1993, Polyachenko 1994). The distinction between them is drawn along a number of lines. Firstly, bars with different rotation velocities have considerably different sizes: while the common bars end up at the corotation or 4:1 resonance, the Lynden-Bell bars end up in the vicinity of the inner Lindblad resonance. Secondly, it is believed that these bars are produced by entirely different formation mechanisms. For the slow Lynden-Bell bars, the physical mechanism is absolutely clear~--- this is the mutual attraction and ``sticking'' together of the slowly precessing orbits. Note, that in application to galactic bars, the idea of axis alignment was first suggested by Lynden-Bell (1979).  Yet for the fast bars the situation remains vague and debatable. The fast bars were long thought to be generated due to the fast rotation in much the same way as in the classical incompressible Maclaurin spheroids being they are strongly oblate and rapidly rotating. It was shown by Toomre (1981) that in reality the galactic bar-modes have little in common with incompressible edge modes. Accordingly, they should have another formation mechanism. Toomre (1981) offered the so-called swing amplification mechanism, which currently is generally accepted for the explanation of the normal SA spirals. We think, however, that a simple extension of the Toomre mechanism to the SB galaxies can hardly be done. In particular, it is difficult to suspect the existence of the running spiral waves, necessary for the swing amplification, which takes place outside the bar near the corotation. The specific attempt made by Athanassoula and Sellwood (1986) to demonstrate the validity of the swing amplification mechanism on some some models of stellar disks appears to be somewhat artificial.  In our opinion, it is more preferably to find such a mechanism of the common bar formation, which is  directly linked to the instability of the central part of the galactic disk itself.  Contopoulos (1975, 1977) has pointed out that the properties of the families of the periodic orbits in the rotating bar potentials play the fundamental role in the theory of the galactic bars. This is especially true for the so-called $x_1$ family of orbits, which is elongated along the bar inside the corotation circle. It is plausible to assume that these periodic and close non-periodic orbits make up the galactic bar. One should note, however, that the preceding is relevant to the theory of the already existing bars, and in the strict sense cannot be applied the the bar formation mechanism itself\\footnote{This is especially true for the linear stage of the bar-forming instability (if the bar arose due to the growth of the instability): the orbits considered by Contopoulos are obviously trapped by the potential of the bar. The trapping process is, of course, the non-linear phenomenon.}.  Nevertheless, from the described picture of orbits constituting the bar, it is customary to make a ``reasonable'' conclusions about the mechanisms of bar formation, including the possible instability at the linear stage. The fast bar angular velocity $\\Omega_p$ is greater (sometimes substantially greater) than the maximum precession velocity $\\Omega_{pr}^{\\max}$. Meanwhile, it is intuitively suggested that the angular modulation of the precessing orbit distribution (i.e., a~figure of the bar) should rotate with the velocity of about mean precession velocity of the orbits. Such a conclusion would have meant the uselessness of the Lynden-Bell mechanism for describing the fast bar formation. So, it is generally agreed that in reality the growing bar forces the orbits to change their shape and tune to the strengthening and narrowing bar. Moreover, it is assumed that this effect can be significant even at the linear stage. It is justified on the example of weak bars in the framework of the linear theory (Sellwood 1993): the theory shows that initially circular orbits change into slightly elongated ovals, with orientation relative to the bar just as the orbits of the $x_1$-family.  However, one can adduce two important objections as to the said in the last paragraph.  1. The circular orbits cannot be the typical sample orbits of the undisturbed disk (excepting for a cold disk model, which is of little interest in this case) until the bar-induced perturbed velocities do not exceed the velocity dispersion in the original axisymmetric disk. It is clear that this condition should be met first at the initial stage of the instability development.  2. The intuitive conclusion mentioned above (that the velocity of the wave density in the system of precessing orbits, $\\Omega_p$, cannot exceed the maximum precession velocity    $\\Omega_{pr}^{\\max}$) is supported by the results of computations of the unstable modes in numerous models of galactic disks. These computations show that in all cases the pattern speed $\\Omega_p$ is smaller than the maximum star angular velocity $\\Omega_{\\max}$ in the disk\\footnote{Note that this property of the oscillation frequency spectrum of the gravitating disk is not yet found for the general case.}. But this analogy appears to be somewhat loosely after the discovering of the possible ``donkey'' behavior of the star orbits by Lynden-Bell and Kalnajs (1972), when the orbit accelerates when held back and slows down when urged forward. Such a behavior takes place under the Lynden-Bell condition, $\\left.\\partial \\Omega_{pr}/\\partial L\\right|_{J_f} < 0$ (Lynden-Bell 1979), where $L$ is the star angular momentum, $J_f = J_r + L/2$ is the Lynden-Bell adiabatic invariant , $J_r$ is the radial action. As it is shown below, if the orbits with a ``donkey'' behavior play an active role in the bar-forming instability, the bar angular velocity can exceed $\\Omega_{pr}^{\\max}$.  The last argument allows to treat the fast bar-modes as the corresponding density waves of the precessing orbits, in the full correspondence with the theory of the Lynden-Bell bars. Note that Kalnajs (1973) was first who called attention to the  possibility of the freely precessing orbit alignment (and thus the bar formation). The important point in his theory of kinematic waves in the case of near-circular orbits was the independence of the precession velocity $\\Omega_{pr}(r) = \\Omega(r) - \\kappa(r)/2$ ($\\Omega(r)$ from radius (here $\\kappa(r)$ is the epicyclic frequency, $\\kappa^2 = 4\\Omega^2 + rd\\Omega^2/dr$). Note that the same condition $\\Omega_{pr}(r) \\approx \\mathrm{const}$ is required in the Lynden-Bell (1979) more general consideration of the problem.  In reality different stars precessing with different velocities. But only important point is that the angular precession velocities (and also the pattern speeds of bars ($\\Omega_p$) in the majority of cases) are substantially smaller than the typical star azimuthal frequencies $\\Omega_2$ and especially the radial frequencies $\\Omega_1$. If for some orbit the inequality \\begin{equation} |\\Omega_p - \\Omega_{pr}|/\\Omega_1 \\ll 1 \\label{cond1} \\end{equation} holds, this whole orbit (not individual stars) can be considered as an elementary object in the interaction with the bar gravitational field. If the condition (\\ref{cond1}) holds for the majority of disk orbits, participating in the bar formation, we can consider the processes (e.g., bar-instability) in the model disk, consisting of the set of the precessing stars. As it is noted by Lynden-Bell (1979), the condition (\\ref{cond1}) means the conservation of the adiabatic invariant $J_f = J_r + L/2$, that reduces the problem of the bar-instability to the one dimensional problem: one have to follow only the variation of the azimuthal positions of the orbit major axes under the action of the gravitational attraction from the bar.  Actually, we solve the problem of bar density wave of orbits with different precession velocities, that is quite analogous to more familiar problem of the star density waves (e.g., of the bar-like shape) in the differentially rotating disks. Seemingly, not all the orbits participating in the bar-mode formation, meet the condition (\\ref{cond1}) with margin, especially for the fast bars. But even for the latter usually $|\\Omega_p - \\Omega_{pr}| \\sim \\Omega_{pr} \\sim \\omega_G$, where $\\omega_G\\sim\\sqrt{GM_d/a^3}$ is the characteristic gravitational (Jeans) frequency ($G$  is the gravitational constant, $M_d$ is the mass of the active disk, $a$ is its radius). Then when $\\Omega_{pr}/\\Omega_1 \\ll 1$ we are still within the bounds of the condition (\\ref{cond1}). Even in case of more weak inequality $\\Omega_{pr}/\\Omega_1 < 1$  (e.g., several times smaller but not at some orders) the suggested model will likely provide the correct qualitative answer. This is at any case not worse than, for example, the analysis of the spiral structure of the galaxy M\\,33 by Shu {\\it et al.} (1971) by using the WKB formulae from the well-known theory of Lin and Shu, applicable, strictly speaking, to the tightly wound multi-turn spirals.  A few words about the content of the paper. In the Section 2, we descibe in the general form a model of the precessing orbits and derive the basic equations for the model, useful for the analysis of the disk low frequency modes of interest for us. The use of the action--angle variables $I_1$, $I_2 = L$ and $w_1$, $w_2$ is the easiest way for derivation of these equations from the general kinetic equation. The substantial simplification appears with introducing a slow angular variable $\\bar w_2 = w_2 - w_1/2$ and averaging over the fast angle variable $w_1$ (i.e., over the radial stellar oscillations). In the simplest cases,  the problem can be reduced to the analysis of a rather simple dispersion relations. In some more general cases, we obtain the integral equations of different degrees of complexity. But even in the most general form, the obtained integral equations for this model is much simpler that the immense integral equations for the normal modes of the stellar disk, obtained by Kalnajs (1965) and Shu (1970). Remind that the use of these general integral equations (not counting the $N$-body methods) was the only possibility for analysis of large-scale modes, as opposed to incomparably simpler problem of tightly-wound spirals . This is especially true in regard to the bar-mode. As discussed above, the real simplification has come through the analysis of low-frequency modes. Note that the resulting description of gravitating systems is similar to the drift approximation in the plasma physics (see, e.g., Chew  et. al., 1956), but for orbits of essentially more general type. In our opinion, the most important advantage of our approach consists in the fact that it makes clear the simple physical mechanisms of the  instability processes developing in a disk. Unfortunately, a possibility of revealing these physical mechanisms under the use of the general integral equations by Kalnajs or Shu would be practically impossible, and the same is true for the $N$-body simulations.  In Sec. 3, we analyze the dispersion relation for a model disk consisting of orbits of two different types  that differ by their precessing speeds ($\\Omega_{pr}^{(1)}$ and $\\Omega_{pr}^{(2)}$; $\\Omega_{pr}^{(2)} > \\Omega_{pr}^{(1)}$) and, generally speaking, also by a sign of the Lynden-Bell derivative  $(\\partial \\Omega_{pr}/\\partial L)_{J_f} \\equiv \\Omega'_{pr}$ .For the case when the derivatives $(\\Omega'_{pr})^{(1)}$ and $(\\Omega'_{pr})^{(2)}$ have opposite signs (i.e., a disk contains both ``attracting'' and ``repelling'' orbits), the pattern speed of the unstable modes ${\\mathrm Re}\\, \\Omega_p$ may be more than $(\\Omega_{pr})^{(2)}$, i.e. the maximum orbit precessing speed in a disk. We show that at the same time $\\Omega_{pr}^{(1)} < \\mathrm{Re}\\,\\Omega_p < \\Omega_{pr}^{(1)}$ when $(\\Omega'_{pr})^{(1)}$ and $(\\Omega'_{pr})^{(2)}$ have the same sign  (positive for instability). We consider this result as the important argument in favor of the unified formation mechanism both for slow and fast bars.  In Sec. 4, the results of computation of the precessing speeds $\\Omega_{pr}$ and the most important for the theory quantity, the Lynden-Bell derivative $\\Omega'_{pr}$, are given for a number of typical potentials. In particular,  the regions on the Lynden-Bell $J_f, L$ plane where $\\Omega'_{pr} > 0$ and $\\Omega'_{pr} < 0$ are found. The natural suggestion is made that a bar forms by ``attracting'' orbits with $\\Omega'_{pr} > 0$\\footnote{Note that in the formation process for all the mode, i.e. both the bar and adjacent spirals, and not just the central bar, the ``repelling'' orbits with $\\Omega'_{pr} < 0$ may also take part. Moreover, this may be important to explain the fast bar phenomenon (as discussed above).}. Then the bar length (more exactly, the radius of the bar's end) should be equal to the maximum of the apogee radii ($r_{\\max}$), among the sufficiently occupied  orbits from the region where $\\Omega'_{pr} > 0$. So, for calculation of $l_b$, one should know not only the general pattern of orbit precessions determined by  the potential $\\Phi_0(r)$ but a specific equilibrium distribution function as well. As a first approximation, one can take the estimate $l_b \\sim r_c$, where $r_c$ is the radius of the circular orbit at which $\\Omega'_{pr} = 0$ ($\\Omega'_{pr} > 0$  for $r < r_c$, and $\\Omega'_{pr} < 0$ for $r > r_c$, at the circular orbits).  These radii $r_c$ are computed for all the potentials considered. Note that our  determination of the bar length has nothing to do with common determinations that link this length with a location of one of the resonances (CR, ILR or 4:1). It is clear  that our bar length may take a great variety of values depending on the specific  potential and distribution function (accidentally, they may fall near one of resonances).  In Sec. 5, we shortly formulate the most important conclusions and discuss some immediate prospects for a work in this field. ",
        "conclusions": "Let us formulate and discuss some conclusions from the theory above.  1. We advanced a number of arguments in favor of universality of the mechanism that may be responsible for formation of both the slow and fast bars. The essence of this mechanism is naturally formulated on the basis of representing a stellar disk as a set of precessing orbits. Such a concept is adequate to the problem under consideration since the bar pattern speeds (including ``fast bars'') are significantly less than the characteristic frequencies of oscillations of individual stars ($\\Omega_1$ and $\\Omega_2$), but they are just of order of the orbit precession speeds ($\\Omega_{pr}$). In such a disk, we seek the unstable normal modes (first of all the bar-mode) as the density wave of precessing orbits that runs with some speed $\\Omega_p$ without any deformations despite the fact that different orbits precess with different speeds (analogously to differentiability of a disk rotation when the usual concept of a galactic disk as a set of individual stars is used).  The unstable bar-mode forms if a central region of a disk (a location of a future bar) contains a sufficiently massive group of ``attracting'' orbits that satisfy the Lynden-Bell condition $\\Omega'_{pr} > 0$, and, besides, the precession speed dispersion of these active orbits are not too large (otherwise, the orbits will run away from the region of perturbation under the influence of the ``thermal'' motion). The last condition is natural for the Jeans nature of the instability under consideration. The exact criteria of instability can be obtained by solution of the basic equations derived in Sec. 2; for some simple cases, the corresponding dispersion relations are given in an explicit form. The pattern speed of a bar $\\Omega_p$ depends significantly on the extent to which the ``repelling'' orbits with $\\Omega'_{pr} < 0$ (the orbits with a ``donkey'' behavior) take part in the bar-formation process. If such orbits are hardly dragged in the bar-formation process, then $\\Omega_p \\sim \\bar\\Omega_{pr}$; just such bars would naturally be named as the ``slow'' bars. But if the role of the ``repelling'' orbits is essential, we can obtain the ``fast'' bars with $\\Omega_p$ that significantly exceeds $\\bar\\Omega_{pr} $(just the same, $\\Omega_p$ should be of order of $\\bar\\Omega_{pr}$ as before --- if we want to remain within the framework of our theory). In Sec. 3, such a possibility is demonstrated on the simplest example of a two-component disk model. Thus, from the point of view under consideration, distinctions between the slow and fast bars are mainly quantitative, but they do not differ fundamentally from each other: both bars form under the action of the same physical mechanism. It is worth noting that the Jeans mechanism (including one under consideration) is always best natural and suited to the gravitational problems.  2. So far the theory above was confirmed only by the calculations of the lowest-frequency modes of the disk models in the Shuster potential when the results are compared with the $N$-body results by Athanassoula \\& Sellwood (1986). These modes obviously correspond to the slow bars: for them, $\\Omega_p \\sim \\bar\\Omega_{pr}$. The most interesting prediction of the theory (its validity for the fast bars) would be tested if we shall be able to prove that the eigen frequencies of the modes computed from our integral equations (see Sec. 2) coincide with the frequencies of the fast bars derived from the $N$-body simulations (in particular, with the majority of bar-modes from the paper by Athanassoula \\& Sellwood cited above. This problem will be the subject of study in the immediate future. In the present paper, we restricted ourselves only to some positive facts that result from the general analysis of orbit precessions for a number of potentials (Sec. 4). These facts correlate with a possibility of the fast bar formation by alignment of ``attracting'' orbits. First of all we noted that the maximum radius of nearly-circular orbits ($r_c$) with $\\Omega'_{pr} > 0$ ($r_c$ may be taken as a natural estimate for a bar end $l_b$) is significantly more than a size ($\\sim r_m$) of the solid rotation region, for all the reasonable models. Moreover, if one takes (as it usually does) that $l_b \\sim r_{CR}$ (where $r_{CR}$ is the corotation radius), then for $l_b \\sim r_c$ one can obtain the estimate $\\omega \\approx 0.46$ for a typical eigen frequency $\\omega$ of bar-modes in the Shuster potential; this estimate is in good agreement with the corresponding results of Athanassoula \\& Sellwood (1986)."
    },
    "0212/astro-ph0212235_arXiv.txt": {
        "abstract": "s{} ",
        "introduction": "\\subsection{The UV-X-$\\gamma$ spectrum of Seyfert galaxies} The typical broad band UV-X-$\\gamma$ ray spectrum of a Seyfert galaxy, here NGC 5548 has been reported in Fig. \\ref{specsey} This spectrum has been obtained from quasi simultaneous observations of different satellites indicated on the Figure \\cite{mag98}. We clearly see different components: \\begin{figure*}[b!] \\vskip -0.7 true cm \\begin{center} \\centerline{\\epsfig{figure=petruccifig1.eps, width=0.8\\textwidth}} \\end{center} \\vskip -0.3 true cm \\caption{Typical broadband UV-X-$\\gamma$ ray spectrum of a Seyfert galaxy.} \\label{specsey} \\end{figure*} \\begin{itemize} \\item A strong emission in the UV, the UV bump, generally interpreted as the thermal emission of an accretion disc. \\item A strong emission in the X-ray band which can be generally fitted by a cut-off power law shape with a power law photon index $\\Gamma\\simeq 1.9$ \\cite{nan94} and a cut-off energy $E_c$ near 200-300 keV \\cite{mat99,per02} \\item A neutral fluorescence Fe$_{K\\alpha}$ line near 6.4 keV \\cite{nan97} \\item A bump peaking near 30 keV \\cite{pou90} \\item A soft component below 2 keV, in excess to the extrapolation of the high energy power law, called the soft excess \\cite{tur89} \\end{itemize}  These different components are generally interpreted in the reprocessing/upscattering model framework. The scheme of this model is reported in Fig. \\ref{repmodel}. It supposes the presence of 2 phases, a cold one, the accretion disc, which produces the UV emission, and a hot one, the corona, a plasma of energetic particles, supposed to radiate X-rays by Compton upscattering the UV photons coming from the disc. Inversely, part of the X-ray emission of the corona is absorbed by the cold material and reprocessed in X-ray. Finally, part of the X-ray emission ($\\sim$ 10\\%) is Compton reflected on the disc surface and give birth to the fluorescent Iron line and the reflection bump. The origin of the soft excess is still unclear, it may be the hard tail of the disc emission.  \\begin{figure*}[h!] \\vskip -0.7 true cm \\begin{center} \\centerline{\\epsfig{figure=petruccifig2.eps, width=8cm}} \\end{center} \\vskip -0.3 true cm \\caption{Schematic view of the reprocessing/upscattering model} \\label{repmodel} \\end{figure*}  The nature of the high energy source(s), the geometry of the disc-corona configuration and the origin of the variability, which is important in Seyfert galaxies, are still unknown. The presence of a high energy cut-off in most of the Seyfert (Matt 2000) and the lack of any annihilation line have favored a thermal nature for the corona even if it does not rule out the non-thermal hypothesis \\cite{pet01}. Recent results on spectral variability (cf. section \\ref{sectnat}) appear to be more conclusive in agreement with the thermal interpretation. We will thus focus below on the thermal solution.  \\subsection{The thermal Comptonization spectrum}\\label{subsec:comp} As said previously, the X-ray emission in Seyfert is believed to be produce by Compton process of UV seed photons on a thermal population of electrons. For a given geometry, the thermal Comptonization spectral shape depends on few parameters i.e. the temperature and optical depth of the corona, $kT_e$ and $\\tau$, and the temperature of the soft photon field $kT_{bb}$.  It generally has a cut-off power law shape (cf. Fig. \\ref{thercompspec}a) and it is often approximate by a spectral function of the type $F(E)\\propto E^{-\\Gamma}\\exp (-E/E_c)$. This may be however a very crude approximation especially in the case of anisotropic geometries. Indeed, the Compton process is relatively sensible to the anisotropy of the seed photon field and the X-ray emission may also be strongly anisotropic. In the case of a slab geometry for example, the X-ray photons of the first Compton scattering order are preferentially emitted backward toward the disk than toward the observer. It thus produces a break (the anisotropy break) in the observed X-ray spectrum as shown in Fig. \\ref{thercompspec}b. In this case, the use of a cut-off power-law shape may lead to erroneous parameter values and to wrong physical interpretations \\cite{pet00,pet01}.  \\begin{figure*}[h] \\begin{center} \\centerline{\\epsfig{figure=petruccifig3b.eps, width=1\\textwidth}} \\end{center} \\caption{Left: A thermal Comptonization spectrum in isotropic geometry. Right: comparison of different Comptonization spectra for different geometry: slab (solid line), hemisphere (dashed line), sphere (i.e. isotropic geometry, dot-dashed line).  The values of $\\tau$ and $kT_e$ in each case have been chosen so that the spectra have roughly the same slope at low energy. We clearly see the anisotropy break present for anisotropic geometries. In dotted line we have plotted a simple cut-off power law approximation.} \\label{thercompspec} \\end{figure*}  It is interesting to note that in the case of radiative equilibrium between the UV (i.e. the cooling) and the X-ray (i.e. the heating) sources, the temperature and optical depth of the corona follow an univocal relationship, which depends however on the geometry of the disc-corona configuration \\cite{haa93,ste95}. This relationship roughly corresponds to a constant Compton parameter $y = \\displaystyle \\left[ 4\\frac{kT_e}{mc^2}+16\\left( \\frac{kT_e}{mc^2}\\right)^2\\right](\\tau+\\tau^2)$. ",
        "conclusions": ""
    },
    "0212/astro-ph0212003_arXiv.txt": {
        "abstract": "Spacetime-varying coupling constants can be associated with violations of local Lorentz invariance and CPT symmetry. An analytical supergravity cosmology with time-varying fine-structure constant provides an explicit example. Estimates are made for some experimental constraints. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/astro-ph0212198_arXiv.txt": {
        "abstract": "We present a detailed study of cosmological effects of homogeneous tachyon matter coexisting with non-relativistic matter and radiation, concentrating on the inverse square potential and the exponential potential for the tachyonic scalar field. A distinguishing feature of these models (compared to other cosmological models) is that the matter density parameter and the density parameter for tachyons remain comparable even in the matter dominated phase. For the exponential potential, the solutions have an accelerating phase, {\\it followed by} a phase with $a(t)\\propto t^{2/3}$ as $t\\to \\infty$. This eliminates the future event horizon present in $\\Lambda CDM$ models and is an attractive feature from the string theory perspective. A comparison with supernova Ia data  shows that for both the potentials there exists a range of models in which the universe undergoes an accelerated expansion at low redshifts and are also consistent with requirements of structure formation. They do require fine tuning of parameters but not any more than in the case of $\\Lambda CDM$ or quintessence models. ",
        "introduction": "Observations suggest that our universe has entered a phase of accelerated expansion  in the recent past. Friedmann equations can be consistent with such an accelerated expansion only if the universe is populated by a medium with negative pressure. One of the possible sources which could provide such a negative pressure will be a scalar field with either of the following two types of Lagrangians: \\begin{eqnarray} L_{\\rm quin} &=& \\frac{1}{2} \\partial_a \\phi \\partial^a \\phi - V(\\phi); \\\\ \\nonumber \\quad L_{\\rm tach} &=& -V(\\phi) [1-\\partial_a\\phi\\partial^a\\phi]^{1/2} \\end{eqnarray} Both these Lagrangians involve one arbitrary function $V(\\phi)$. The first one $L_{\\rm quin}$,  which is a natural generalization of the Lagrangian for a nonrelativistic particle, $L=(1/2)\\dot q^2 -V(q)$, is usually called quintessence (for a sample of models, see Ref. \\cite{phiindustry}). When it acts as a source in Friedmann universe, it is characterized by a time dependent   $w(t)\\equiv (P/\\rho)$  with \\begin{eqnarray} \\rho_q(t) &=& \\frac{1}{2} \\dot\\phi^2 + V; \\quad P_q(t) = \\frac{1}{2}; \\dot\\phi^2 - V; \\\\ \\nonumber w_q &=& \\frac{1-(2V/\\dot\\phi^2)}{1+ (2V/\\dot\\phi^2)}. \\label{quintdetail} \\end{eqnarray}  Just as $L_{\\rm quin}$ generalizes the Lagrangian for the nonrelativistic particle, $L_{\\rm tach}$ generalizes the Lagrangian for the relativistic particle \\cite{tptirth}. A relativistic particle with a (one dimensional) position $q(t)$ and mass $m$ is described by the Lagrangian $L = -m \\sqrt{1-\\dot q^2}$. It has the energy $E = m/  \\sqrt{1-\\dot q^2}$ and momentum $p = m \\dot q/\\sqrt{1-\\dot q^2} $ which are related by $E^2 = p^2 + m^2$.  As is well known, this allows the possibility of having massless particles with finite energy for which $E^2=p^2$. This is achieved by taking the limit of $m \\to 0$ and $\\dot q \\to 1$, while keeping the ratio in $E = m/  \\sqrt{1-\\dot q^2}$ finite.  The momentum acquires a life of its own,  unconnected with the velocity  $\\dot q$, and the energy is expressed in terms of the  momentum (rather than in terms of $\\dot q$)  in the Hamiltonian formulation. We can now construct a field theory by upgrading $q(t)$ to a field $\\phi$. Relativistic invariance now  requires $\\phi $ to depend on both space and time [$\\phi = \\phi(t, {\\bf x})$] and $\\dot q^2$ to be replaced by $\\partial_i \\phi \\partial^i \\phi$. It is also possible now to treat the mass parameter $m$ as a function of $\\phi$, say, $V(\\phi)$ thereby obtaining a field theoretic Lagrangian $L =- V(\\phi) \\sqrt{1 - \\partial^i \\phi \\partial_i \\phi}$. The Hamiltonian  structure of this theory is algebraically very similar to the special  relativistic example  we started with. In particular, the theory allows solutions in which $V\\to 0$, $\\partial_i\\phi \\partial^i\\phi \\to 1$ simultaneously, keeping the energy (density) finite.  Such solutions will have finite momentum density (analogous to a massless particle with finite  momentum $p$) and energy density. Since the solutions can now depend on both space and time (unlike the special relativistic example in which $q$ depended only on time), the momentum density can be an arbitrary function of the spatial coordinate. This form of scalar field arises  in string theories \\cite{asen} and --- for technical reasons --- is called a tachyonic scalar field. (The structure of this Lagrangian is similar to those analyzed in a wide class of models called {\\it K-essence}; see for example, Ref. \\cite{armen}.) This provides a rich gamut of possibilities in the context of cosmology \\cite{tptirth,armen,cosmos,earl,tptachyon}.   The stress tensor for the tachyonic scalar  field can be written in a perfect fluid form \\begin{equation} T^i_k = (\\rho + p) u^i u_k - p \\delta^i_k \\end{equation} with \\begin{eqnarray} u_k &=& \\frac{\\partial_k \\phi}{\\sqrt{\\partial^i \\phi \\partial_i \\phi}};\\quad \\rho = \\frac{V(\\phi)}{\\sqrt{1 - \\partial^i \\phi \\partial_i \\phi}};\\quad  \\\\ \\nonumber p &=& -V(\\phi) \\sqrt{1 - \\partial^i \\phi \\partial_i \\phi} \\end{eqnarray} The remarkable feature of this stress tensor is that it could be considered as {\\it the sum of a pressure less dust component and a cosmological constant} \\cite{tptirth}. To show this explicitly, we  break up the density $\\rho$ and the pressure $p$ and write them in a more suggestive form as \\begin{equation} \\rho = \\rho_V  + \\rho_{\\rm DM}; ~~ p = p_V  + p_{\\rm DM}, \\end{equation} \\\\ where \\begin{eqnarray} \\rho_{\\rm DM} &=& \\frac{V(\\phi) \\partial^i \\phi \\partial_i \\phi} {\\sqrt{1 - \\partial^i \\phi \\partial_i \\phi}};\\ \\  p_{\\rm DM} = 0; \\\\ \\nonumber \\rho_V  &=& V(\\phi) \\sqrt{1 - \\partial^i \\phi \\partial_i \\phi};\\ \\ p_V  = -\\rho_V \\end{eqnarray} This means that the stress tensor can be thought of as made up of two components -- one behaving like a pressure-less fluid, while the other having a negative pressure. In the cosmological context, when $\\dot\\phi$ is small (compared to $V$ in the case of quintessence or compared to unity in the case of tachyonic field), both these sources have $w\\to -1$ and mimic a cosmological constant. When $\\dot \\phi \\gg V$, the quintessence has $w\\approx 1$ leading to    $\\rho_q\\propto (1+z)^6$; the tachyonic field, on the other hand, has $w\\approx 0$ for $\\dot\\phi\\to 1$    and behaves like non-relativistic matter.  An additional motivation for studying models based on $L_{\\rm quin}$ is the following: The standard explanation of the current cosmological observations will require two components of dark matter: (a) The first one is a dust component with the  equation of state  $p=0$ contributing $\\Omega_m \\approx 0.35$. This component clusters gravitationally at small scales  ($l \\lesssim 500$ Mpc, say) and will be able to explain observations from galactic to super-cluster scales. (b) The second one is a negative pressure component with equation of state like $p=w\\rho$  with $-1 < w < -0.5 $ contributing about $\\Omega_V \\approx 0.65$. There is  some leeway in the $(p/\\rho)$ of the second component but it is certain that  $p$ is negative and $(p/\\rho)$ is of order unity. The cosmological constant will provide $w=-1$ while several other candidates based on scalar fields with potentials \\cite{phiindustry} will provide different values for $w$ in the acceptable range.  By and large, component (b) is noticed only in the large scale expansion and it does not cluster gravitationally to a significant extent. Neither of the components (a) and (b) has laboratory evidence for their existence directly or indirectly. In this sense, cosmology requires invoking the tooth fairy twice to explain the current observations. It was suggested recently \\cite{tptirth} that one may be able to explain the observations at all scales using a single scalar field with a particular form of Lagrangian.  In this paper we explore the cosmological scenario in greater detail, concentrating on the background cosmology. Our approach in this paper will be based on the above view point and we shall treat the form of the Lagrangian for $L_{\\rm tach}$ as our starting point without worrying about its    origin. In particular, we shall {\\it not}     make any attempt to connect up the form of $V(\\phi)$ with string theoretic models but will explore different possibilities,    guided essentially by their cosmological viability.   We construct cosmological models with homogeneous tachyon matter, assuming that tachyon matter co-exists with normal non-relativistic matter and radiation. Section \\ref{sec:tachmatter} presents the equations for the evolution of scale factor and the tachyon field. In Section \\ref{sec:backgd_solns} we present solutions of these equations and discuss variations introduced by the available parameters in the model. Here we have analyzed two different models of the scalar field potential, one is the exponential potential and the other is the the inverse square potential which leads to power law cosmology \\cite{tptachyon}. In Section \\ref{sec:supnova_data} we compare the model with observations and constrain the parameters. Section \\ref{sec:tachstr} discusses structure formation in tachyon models. The results are summarized in concluding Section \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have shown that it is possible to construct viable models with tachyons contributing significantly to the energy density of the universe.  In these models, matter, radiation and tachyons co-exist. We show that a subset of these models satisfy the constraints on the accelerating expansion of the universe. For the accelerating phase to occur at the present epoch, it is necessary to fine tune the initial conditions.  We have further demonstrated that the density parameter for tachyons does not becomes negligible at high redshifts, hence the growth of perturbations in non-relativistic matter is slower for most models than, e.g., the  $\\Lambda$CDM model. This problem does not effect a small subset of models. However, given that the density parameter of tachyons cannot be ignored in the ``matter dominated era'', it is essential to study the fate of fluctuations in the tachyon field."
    },
    "0212/astro-ph0212151_arXiv.txt": {
        "abstract": "We present results from a two week multi-longitude photometric campaign on TV~Col held in 2001 January. The data confirm the presence of a permanent positive superhump found in re-examination of extensive archive photometric data of TV~Col. The 6.3-h period is 15 per cent longer than the orbital period and obeys the well known relation between superhump period excess and binary period. At 5.5-h, TV~Col has an orbital period longer than any known superhumping cataclysmic variable and, therefore, a mass ratio which might be outside the range at which superhumps can occur according to the current theory. We suggest several solutions for this problem. ",
        "introduction": "\\subsection{Permanent superhumps}  Patterson \\& Richman (1991) initially suggested the term `permanent superhump' for the subclass of cataclysmic variables (CVs) having quasi-periodicities slightly different from their binary orbital periods. Unlike SU~UMa systems (see Warner 1995 for a review of SU~UMa systems and CVs in general), which show this behaviour only during superoutbursts, permanent superhump systems show the phenomenon during their normal brightness state.  However, their amplitudes are highly variable and are sometimes below the detection limits, so the term `permanent superhump' is somewhat misleading.  Whitehurst \\& King (1991) suggested that superhumps occur when the accretion disc extends beyond the 3:1 resonance radius.  According to Osaki (1996), permanent superhumpers differ from other subclasses of non-magnetic CVs in having relatively short orbital periods and high mass-transfer rates, resulting in accretion discs that are thermally stable but tidally unstable.  Retter \\& Naylor (2000) provided observational support for this idea.  The `positive superhump', a periodicity that is a few per cent larger than the orbital period, is explained as the beat period between the binary motion and the precession of an eccentric accretion disc in the apsidal plane. Periods slightly shorter than the orbital periods have also been seen in several systems.  They are known as `negative superhumps', and may be generated by the nodal precession of the accretion disc (Patterson et al. 1993; Patterson 1999).  However, there are some theoretical difficulties with this idea (Murray \\& Armitage 1998; Wood, Montgomery \\& Simpson 2000; Murray et al. 2002).  Superhumps were also associated with the formation of a spiral structure in the accretion disc (Steeghs, Harlaftis \\& Horne 1997; Baba et al. 2002).  Observations of positive superhumps have shown a roughly linear relationship between the period excess, expressed as a fraction of the binary period, and the binary period itself (Stolz \\& Schoembs 1984). Negative superhumps seem to obey a similar rule (Patterson 1999).      \\begin{table*} \\begin{minipage}{300 mm}  \\caption{\\label{table.all} Photometric observations of TV Col} \\begin{tabular}{@{}lccccccccc@{}}  \\hline  Set & Month/year & Nights & Site     & Telescope size  & Detector   & Filter/s  & Exposures & Mean night & Number of \\\\ &            &        &           & [m]            &            &           & [s]      & length [h] & outbursts \\\\ \\hline  1   & 12/1985      & 6      & ESO      & 0.91           & photometer & V,B,L,U,W & 16        & 4.2 (gaps) & 0       \\\\  2   &12/1985--1/1986  & 13     & SA       & 0.75+1.0       & photometer & white     & 2        & 3.0        & 0      \\\\  3   &11--12/1987    & 16     & ESO      & 0.91           & photometer & V,B,L,U,W & 16        & 5.4 (gaps)  & 3      \\\\  4   & 12/1987      & 9      & ESO      & 0.91           & photometer & V,B,L,U,W & 16        & 5.0        & 0       \\\\  5   & 11/1988      & 7      & ESO      & 0.91           & photometer & V,B,L,U,W & 16        & 4.4 (gap)  & 0       \\\\  6   & 1/1989       & 6      & SA       & 1.0            & CCD        & white     & 40        & 7.3        & 0      \\\\  7   & 1/1991       & 9      & SA       & 0.75           & photometer & B,R       & 4         & 4.8        & 0      \\\\  8   & 12/1991      & 11     & SA       & 0.75           & photometer & white     & 10        & 4.0        & 1      \\\\  9   & 1/2001       & 5      & SA+AU+NZ & 0.75+0.40+0.25 & CCD        & white     & 20-90     & 11.8       & 0      \\\\  10  & 1/2001       & 8      & SA+AU+NZ  & 0.75+0.40+0.25 & CCD        & white     & 20-90     & 10.6       & 0      \\\\   \\hline  \\end{tabular} \\end{minipage} \\end{table*}  \\subsection{TV~Col}   TV~Col is a 14th mag CV at a distance of about 370 pc (McArthur et al.\\ 2001). Its light curve shows multiple periodicities that have attracted many observers. Motch (1981) found periodicities of 5.2\\,h and 4\\,d from photometry taken over 10 nights.  A radial velocity study by Hutchings et al.\\ (1981) confirmed these periods and detected another at 5.5 h, which they identified as the orbital period.  They also pointed out that the 4-d period is exactly the beat periodicity between the other two periods. The 5.2-h period was interpreted as the spin period of a magnetic white dwarf, leading to a classification of the object as an intermediate polar (for reviews of intermediate polars, see Patterson 1994; Warner 1996; Hellier 1996). The detection of a 32-m period in x-ray observations (Schrijver et al.\\ 1985; Schrijver, Brinkman \\& van der Woerd 1987) confirmed the intermediate polar nature of TV~Col, but left the 5.2-h period unidentified.  Further observations confirmed the presence of the three periods in the optical regime (Barrett, O'Donoghue \\& Warner 1988; Hellier, Mason \\& Mittaz 1991; Hellier 1993; Augusteijn et al.\\ 1994). The orbital period was also detected in the ultraviolet (Bonnet-Bidaud, Motch \\& Mouchet 1985). It was found that the 5.2-h period is not stable (e.g. Augusteijn et al.\\ 1994), suggesting that it is a superhump period. The fact that it is shorter than the binary period would make it a negative superhump.  Barrett et al.\\ (1988), however, discussed other possible models for the system.  \\subsection{The significance of superhumps in TV~Col}  Superhumps have only been observed in CVs with orbital periods below about 3.7\\,h.  Since the orbital period of TV~Col is 5.5 h, the interpretation of its 5.2-h period as a negative superhump would make this system an extreme case.  There is known to be a strong correlation in CVs between the orbital period and the mass of the secondary star (Smith \\& Dhillon 1998), which arises because systems with longer orbital periods have larger separations and thus require more massive secondaries to fill their Roche-lobes.  The 5.5-h orbital period of TV~Col implies a secondary star of $\\sim 0.6 M_{\\odot}$ -- much more massive than in other superhumping systems. A typical white dwarf in CVs has a mass of $\\sim 1.0 M_{\\odot}$, which means the mass ratio, $q=M_{\\rm donor}/M_{\\rm compact}$, in TV~Col could be about 0.6 or even higher.  This value significantly exceeds the theoretical limit for superhumps -- $q\\la$0.33 (Whitehurst 1988; Whitehurst \\& King 1991; Murray 2000). TV~Col thus offers a unique opportunity to test the predictions of the models for precessing accretion discs.  There seems to be a strong connection between positive and negative superhumps. Light curves of many permanent superhumpers show both types of superhumps (Patterson 1999; Arenas et al.\\ 2000; Retter et al.\\ 2002). In addition, period deficits in negative superhumps are about half the period excesses in positive superhumps (Patterson 1999): $\\epsilon_{\\rm negative}\\approx -0.5 \\epsilon_{\\rm positive}$, where $\\epsilon =(P_{\\rm superhump} - P_{\\rm orbital})/P_{\\rm orbital}$.  This prompted us to examine available photometry of TV~Col for positive superhumps, which would be predicted to have a period near 6.3\\,h.  We indeed found such evidence (Retter \\& Hellier 2000; Retter et al.\\ 2001), which is presented here.  At the request of a referee, we then obtained new multi-site photometry on TV~Col in 2001 January, which confirmed the 6.3-h period. ",
        "conclusions": "\\begin{enumerate}  \\item We have found a 6.3-h period in the optical light curve of TV~Col in the multi-longitude photometric campaign held in 2001 January.  This detection confirms our findings from re-analysis of previously published photometric data.  The periodicity is most naturally explained as a positive superhump.  \\item In the middle of the 2001 January run a short-lived minor outburst occurred.  The new period is seen only in the data following outburst. The outburst also changed the phase of the negative superhump by about a quarter of a cycle, while its amplitude was slightly decreased.  \\item Our findings support the classification of TV~Col as a permanent superhump system.  Its superhump period is the longest known among CVs. TV~Col thus offers a unique opportunity to test and reject some of the models, as it extends the superhump regime to periods far beyond the predicted values into a regime where the differences between models become significant.  The mass ratio of TV~Col might exceed the limit for superhump systems allowed by hydrodynamic simulations.  \\item The observational and physical link between positive and negative superhumps is thus strengthened by our result.  We might speculate that all permanent superhump systems may have both types of superhumps.     \\item Our finding significantly extends the upper limit of orbital periods in positive superhump systems (from 3.7 h to 5.5 h).  This range of periods is populated with dwarf novae and nova-like systems.  Therefore, we strongly urge observers to search for superhumps in nova-likes and dwarf novae with orbital periods above the period gap and up to at least 5.5 h, to check whether TV~Col is unique.  If TV~Col has certain properties that can explain its extraordinarily long superhump period (e.g. classification as an intermediate polar), then other long-period superhump systems should possess similar features.   \\end{enumerate}"
    },
    "0212/astro-ph0212367_arXiv.txt": {
        "abstract": "We find a new mechanism of neutron star radiation wherein radiation is produced by the stellar interior. The main finding is that neutron star interior is transparent for collisionless electron sound, the same way as it is transparent for neutrinos. In the presence of magnetic field the electron sound is coupled with electromagnetic radiation, such collective excitation is known as a fast magnetosonic wave. At high densities the wave reduces to the zero sound in electron liquid, while near the stellar surface it is similar to electromagnetic wave in a medium. We find that zero sound is generated by superfluid vortices in the stellar core. Thermally excited helical vortex waves produce fast magnetosonic waves in the stellar crust which propagate toward the surface and transform into outgoing electromagnetic radiation. The magnetosonic waves are partially absorbed in a thin layer below the surface. The absorption is highly anisotropic, it is smaller for waves propagating closer to the magnetic field direction. As a result, the vortex radiation is pulsed with the period of star rotation. The vortex radiation has the spectral index $ \\alpha \\approx -0.45$ and can explain nonthermal radiation of middle-aged pulsars observed in the infrared, optical and hard X-ray bands. The radiation is produced in the star interior, rather then in magnetosphere, which allows direct determination of the core temperature. Comparing the theory with available spectra observations we find that the core temperature of the Vela pulsar is $T\\approx 8\\times 10^8$K, while the core temperature of PSR B0656+14 and Geminga exceeds $2\\times 10^8$K. This is the first measurement of the temperature of a neutron star core. The temperature estimate rules out equation of states incorporating Bose condensations of pions or kaons and quark matter in these objects. The estimate also allows us to determine the critical temperature of triplet neutron superfluidity in the Vela core $T_c=(7.5\\pm 1.5)\\times 10^9$K which agrees well with recent data on behavior of nucleon interactions at high energies. We also find that in the middle aged neutron stars the vortex radiation, rather then thermal conductivity, is the main mechanism of heat transfer from the stellar core to the surface. The core radiation opens a possibility to study composition of neutron star crust by detection absorption lines corresponding to the low energy excitations of crust nuclei. Bottom layers of the crust may contain exotic nuclei with the mass number up to 600 and the core radiation creates a perspective to study their properties. In principle, zero sound can also be emitted by other mechanisms, rather than vortices. In this case the spectrum of stellar radiation would contain features corresponding to such processes. As a result, zero sound opens a perspective of direct spectroscopic study of superdense matter in the neutron star interiors. ",
        "introduction": "Properties of matter at densities much larger than nuclear are poorly known and constitute a challenging problem of modern science. They have broad implications of great importance for cosmology, the early universe, its evolution, for compact stars and for laboratory physics of high-energy nuclear collisions. Present searches of matter properties at high density take place in several arenas. One of them is investigation of neutron stars.  According to theory of neutron stars (NSs) \\cite{Yako99a} a NS can be subdivided into the atmosphere and four internal regions: the outer crust, the inner crust, the outer core and the inner core as shown in Fig.~1. The atmosphere is a thin plasma layer, where thermal electromagnetic radiation is formed. The depth of the atmosphere varies from some ten centimeters in a hot NS down to some millimeters in a cold one. The outer crust extends from the bottom of the atmosphere to the layer of density $\\approx 4.3\\times 10^{11}{\\rm g}$/${\\rm cm}^3$ and has a depth of several hundred meters. Ions and electrons compose its matter. The depth of the inner crust may be as large as several kilometers. The inner crust consists of electrons, free neutrons and neutron-rich atomic nuclei. Neutrons in the inner crust are in a superfluid state. The outer core occupies the density range $0.86-2\\psi _0$% , where $\\psi _0=2.8\\times 10^{14}{\\rm g}$/${\\rm cm}^3$ is the saturation nuclear matter density, and can be several kilometers in depth. It consists of neutrons $n$ with some (several per cent by particle number) admixture of protons $p$ and electrons $e$. Almost all theories of NSs predict the appearance of neutron superfluidity and proton superconductivity in the outer NS core. In a low-mass NS, the outer core extends to the stellar center.  More dense NSs also possess an inner core (see Fig.~1). Its radius may reach several kilometers and central density may be as high as $10-15\\psi _0$. The composition and equation of state of the inner core are poorly known and it constitutes the main NS ``mystery''. Several hypotheses are discussed in the literature. One of them is appearance of $\\Sigma $- and $\\Lambda $-hyperons in the inner core. The second hypothesis assumes the appearance of pion condensation. The third hypothesis predicts a phase transition to strange quark matter composed of almost free $u$, $d$ and $s$ quarks with a small admixture of electrons. Several authors also considered the hypothesis of kaon condensation in the dense matter. The radius of the NS usually decreases as the stellar mass increases (NS becomes more compact). The mass reaches its maximum at some density, which corresponds to the most compact stable stellar configuration. Typically the stellar mass ranges between $% 0.5-3M_{\\odot }$ for different equations of state ($M_{\\odot }$ is the solar mass).  Despite of 30 years enormous theoretical and observational efforts, the properties of matter in neutron star cores remain poorly known. The main reason is that detected thermal radiation from the star surface carries practically no information about the stellar interior. Some indirect methods, such as estimates of neutron star masses, radii and cooling rates for different equation of states, have been applied to find constrains on the matter behavior. However, non of them is reliable enough because the answer strongly depends on the model parameters which are also poorly known.  In this paper we discuss a new direct way to determine properties of the superdense matter. We found that neutron star interior is transparent for electron zero sound, the same way as it is transparent for neutrinos. Zero sound is a collective density (or spin-density) wave in Fermi liquids that was predicted by Landau in 1957 and later observed in $^3$He and electron liquid in metals. The wave frequency $\\omega $ is greater then the frequency of interparticle collisions which makes the wave different from usual hydrodynamic sound. In neutron stars the zero sound can propagate in electron fluid which fills the stellar crust and core, the collisionless (zero sound) regime is reached at $\\omega /2\\pi >10^{12}$Hz. The large density of electrons in NS interior results in substantial reduction of the interaction energy between fermions as compared to their kinetic (Fermi) energy, the electron gas becomes more ideal. In such weakly interacting Fermi system the attenuation of zero sound is exponentially small (see Eq. (% \\ref{at8}) in Appendix A). As a result, zero sound propagates across the star in a ballistic regime, which is similar to neutrino. However, this is not the case for single-particle excitations in electron liquid, their mean free path usually does not exceed $0.01-1{\\rm cm}$. The single-particle excitations are in a local thermodynamic equilibrium with the stellar matter, while collective zero sound modes do not. At the same time, negligible attenuation of the zero sound creates an obstacle to excite such waves inside the stellar medium. Some special processes are required for their generation. We found that zero sound is generated at the core-crust interface by helical motion of superfluid vortices. Vortex modes themselves are excited thermally via single-particle channel, they are in thermodynamic equilibrium with the surrounding medium.  In the presence of magnetic field the electron sound is coupled with electromagnetic radiation, such excitation is known as a fast magnetosonic wave. At large densities the effect of magnetic field is negligible and the wave behaves as a sound wave, however near the stellar surface (small matter density) it reduces to electromagnetic wave which leave the star and propagates into surrounding space. Such property results in a new mechanism of star radiation produced at the stellar interior. Contrary to neutrino, electromagnetic radiation is easy to detect which opens a perspective of direct spectroscopic study of superdense matter. Thermally excited helical vortex waves generate fast magnetosonic waves in the stellar crust which propagate toward the surface and transform into outgoing electromagnetic radiation. We found that such mechanism results in radiation with non Planck's spectrum (Eqs. (\\ref{mss}), (\\ref{m104})) and theoretical predictions agree well with nonthermal radiation of some middle-aged pulsars observed in the infrared, optical and hard $X$-ray bands (Figs. 5, 6). The important result of our theory is that detection of vortex radiation in the $% X-$ray band allows direct determination of the core temperature. From available spectral data we found that the core temperature of the Vela pulsar is $T\\approx 8\\times 10^8$K, while the core temperature of PSR B0656+14 and Geminga exceeds $2\\times 10^8$K. This is the first measurement of the temperature of a NS core. The cooling rate and temperature of the central part of a NS substantially depends on the properties of dense matter. Our temperature estimate excludes exotic equation of states incorporating Bose condensations of pions or kaons and quark matter in these objects. Such conclusion agrees with recent measurement of gravitational redshift of a neutron star \\cite{Cott02}. Moreover, the temperature estimate allows us to determine the temperature of superfluid transition of neutrons in the Vela core $T_c=(7.5\\pm 1.5)\\times 10^9$K. This value agrees well with theoretical predictions that are based on new nucleon-nucleon interaction potentials which fit the recent world data for $pp$ and $np$ scattering up to $500{\\rm Mev}$ \\cite{Bald98,Arnd97}. Examples of such potentials are Nijmegen I and II. For the matter density $2\\psi _0$ (which is a central region density of a canonical NS without exotic inner core) these nucleon potentials predict $T_c\\approx 7.5\\times 10^9$K for neutron superfluidity with $^3P_2$ (triplet) pairing \\cite{Bald98}. One should mention that in order to find the transition temperature at densities up to $10\\psi _0$ more accurate potentials which fit the measured nucleon-nucleon phase shifts up to about $1{\\rm Gev}$ are required. A systematic study of pulsar radiation in the $X-$ray band can locate objects with fast cooling core which could be candidates for NSs with exotic states of matter.  The plan of the paper is the following. In section 2 we discuss neutron vortices in the superfluid core and generation of magnetosonic waves by vortices at the crust-core interface. In sections 3, 4 we investigate propagation of magnetosonic waves across the stellar crust and their transformation into electromagnetic radiation at the surface. In section 5 we consider spectrum of electromagnetic radiation produced by vortices and analyze the applicability of our results. In section 6 we compare the theory with radiation spectrum of some middle-aged pulsars observed in infrared$-X$% -ray bands and discuss properties of superdense matter based on the observed spectra. Section 7 is devoted to a problem of heat transport across the stellar crust. In section 8 we draw our conclusions and discuss perspectives in the field. In Appendices A, B, C we consider zero sound in an electron liquid, derive equations of collisionless magnetic hydrodynamics for electron motion in the stellar crust and investigate propagation of electromagnetic waves across the atmosphere. ",
        "conclusions": "In this paper we discuss a new possible mechanism of NS radiation produced by vortices in the superfluid core. The key idea is that neutron star interior is transparent for electron zero sound, the same way as it is transparent for neutrinos. The waves of zero sound (in electron liquid) are generated at the core-crust interface by superfluid vortices which undergo thermal helical oscillations. Sound waves produced by vortices propagate across the stellar crust towards the surface. In the presence of magnetic field sound is coupled with electromagnetic waves. Virtually, they are limiting cases of the same essence - the fast magnetosonic wave which is known in magnetic hydrodynamics. Near the crust bottom, where the density is large, the fast magnetosonic waves reduce to zero sound (the effect of magnetic field is negligible). However, near the stellar surface the Alfven velocity becomes larger than the speed of sound and the fast magnetosonic waves behave as electromagnetic waves in a medium. As a result, at the stellar surface the fast magnetosonic waves transform into electromagnetic radiation in vacuum the same way as refraction of usual electromagnetic waves at the dielectric-vacuum interface. For typical magnetic fields of the middle-aged NSs the transformation coefficient is close to $1$.  Finite conductivity of the crust results in partial absorption of magnetosonic waves in a thin layer near the stellar surface. As a result, only a fraction of waves which in the absorbing layer propagate in the direction close to magnetic lines reach the surface and transform into electromagnetic radiation. The fraction depends on the magnetic field strength and surface temperature. For typical parameters of middle aged NSs the fraction is less than $1\\%$, however for colder stars, such as Geminga, it can be substantially larger (even of the order of $100\\%$). As a consequence of anisotropic attenuation and star rotation the observed vortex radiation is pulsed, pulse shape is broader for colder stars. The part of vortex radiation absorbed near the surface contributes to the heat transfer from the core to the surface. Estimates show that at least for some NSs such mechanism, rather than thermoconductivity, is the main source of heating the stellar surface.  Stellar radiation produced by vortices has the spectral index $\\alpha \\approx -0.45$ and for middle-aged pulsars dominates in the IR, optical and hard $X-$ray bands. While the thermal surface radiation contributes mainly in the ultraviolet and soft $X-$ray bands. The observed spectra of PSR B0656+14, Vela and Geminga pulsars agree with our theory. In particular, there is an excellent agreement with observations in two predictions where the theory contains no free parameters: the spectral index of vortex radiation $\\alpha $ and the relation between the core and surface temperatures (given by Eq. (\\ref{d15})). There are also several predictions which have to be confirmed by future observations. For example, the theory predicts a kink of the vortex radiation spectrum in the infrared band (see Eqs. (\\ref{m102}), (\\ref{m104}) and Figs. 5,6) and an exponential decay of the spectrum in the hard $X-$ray band at $\\hbar \\omega >k_BT$ (see Eq. (\\ref {mss})) which was not measured in details yet.  The vortex radiation is produced in the stellar interior and contains important information about properties of the NS core, in particular, it allows direct determination of the core temperature. Comparing theory with available spectra observations we found that the core temperature of the Vela pulsar is $T\\approx 8\\times 10^8$K, which is the first measurement of the temperature of a neutron star core. It is important to note that the temperature we found is much larger than those predicted in heat transport models discussed in literature so far. To understand the big difference we consider heat transport in a model with magnetic field parallel to the core surface (which is just a boundary condition at the superconductor boundary) and show that the effect of magnetic thermal isolation can explain so high value of the core temperature as compared with the stellar surface temperature.  A direct measurement of the core temperature by means of detection the vortex radiation opens a new perspective in study superdense matter. In particular, our estimate of the interior temperatures of PSR B0656+14, Vela and Geminga pulsars rules out possibility of exotic states, such as Bose condensation of pions or kaons and quark matter, in these objects. Also it allows us to extract the superfluid transition temperature for neutrons in the center region of the stellar core $T_c=(7.5\\pm 1.5)\\times 10^9$K which is the first determination of such parameter from observations.  Detection of vortex radiation opens a possibility to study composition of NS crust (NS interior spectroscopy). Since magnetosonic waves generated by vortices propagate through the stellar interior the spectrum of vortex radiation should contain (redshifted) absorption lines which correspond to low energy excitations of nuclei that form NS crust. E. g., the $_{26}^{57}$% Fe nucleus has an excited state with the energy $14.4$kev$=3.5\\times 10^{18}$% Hz, which would produce an absorption line in the hard $X-$ray band of vortex radiation spectrum if temperature of the stellar core is greater than $1.7\\times 10^8$K. The other possible absorption lines in the $X$ ray band can be produced by nuclei $_{21}^{45}$Sc ($12.4$kev), $_{25}^{56}$Mn ($26$% kev), $_{32}^{73}$Ge ($13.5$kev) or by heavy nuclei $_{92}^{235}$U ($13$% kev), $_{97}^{249}$Bk ($8.8$kev). One should mention that bottom layers of neutron star crust may contain exotic nuclei with mass number up to about $% 600$ and vortex radiation can be a possible tool to study their properties. However, the absorption lines of nuclei are narrow and large spectral resolution is necessary for their detection. If temperature of the stellar crust $T_{crust}\\sim 10^8$K the relative thermal line width is $\\sqrt{% 2k_BT_{crust}/Mc^2}\\approx 2\\times 10^{-5}$. Under such condition the star rotation gives the main contribution to the line broadening and the corresponding Doppler line width for middle-aged pulsars can be estimated as $\\Omega R_s/c=0.001-0.01$.  Another challenging goal for future study is interaction of zero sound with exotic states of matter which might exist in more dense NSs. The point is that generation of zero sound by vortices is only one of the possible mechanisms. If an exotic state of matter absorbs zero sound at some frequency, then, according to Kirchhoff's law, it will radiate zero sound at this frequency. As a result, spectrum of stellar radiation must contain characteristic emission lines corresponding to such processes which allows a direct spectroscopic study of superdense matter. In the limiting case when the inner core matter absorbs all zero sound waves the stellar radiation spectrum should be close to radiation of a black body with the temperature of the core. One should look for such objects among bright galactic $X-$ray sources. Also zero sound can be emitted in reactions between elementary particles that constitute the stellar core. It is known that some pulsars are powerful sources of $\\gamma -$rays and it would be interesting to study the connection between $\\gamma $ radiation and zero sound. It is also interesting to investigate a possible connection between electron sound and nonthermal radiation of young (Crab like) pulsars which reveals flat spectrum in the optical band."
    },
    "0212/astro-ph0212017_arXiv.txt": {
        "abstract": "\\psr is an X-ray emitting radio pulsar which was observed with Advanced CCD Imaging Spectrometer (ACIS) on board \\cha on 2001 November 04 and on 2002 March 25 for 2000 s and 17000 s, respectively. The source countrate was 0.0111$ \\pm$0.0013cps. Making \\psr to be the longest period X-ray emitting pulsar. The spectral distribution of counts can be described by several model fits. A blackbody fit yields a temperature $kT=2.69_{-0.09}^{+0.09}\\times 10^{6}$K together with $N_{H}=0.082_{-0.03}^{+0.04}\\times 10^{22}$cm$^{-2}$ and a powerlaw fit yields a photon index of $\\gamma =2.65_{-0.29}^{+0.28}$ with a hydrogen column density of $N_{H}=0.13_{-0.021}^{+0.028}\\times 10^{22}$cm$^{-2}$. Confirming the previous \\ros pointed observation for PSR 1813-36, there was no positive detection from the 30ks \\cha ACIS observation on 2001 October 25. We obtained an upper limit for the countrate of $2.3\\times 10^{-4}$cps. ",
        "introduction": "Radio channel is the best way to discover neutron stars, but their characteristics are best revealed via the X-ray observations i.e. whether it's a cooling neutron star or not. With this in mind, we observed with \\cha two \\ros All-Sky-Survey (RASS) pulsars, that were claimed to be detected, to confirm their detection and examine their X-ray pulses and spectra. During RASS PSR 1813-36 and \\psr had detection likelihood just below the detection threshold, Becker et al. (1992). Therefore both pulsars had follow-up \\ros pointed observations. PSR 1813-36 was not detected and we also report a non-detection of the source for the \\cha observation. However \\psr was detected with $\\sim$60 counts in the \\ros pointed observation and here we report a positive detection of a point source around the pulsar's radio position. Due to low countrate spectral analysis did not suffice to distinguish between various pulsar models such as, powerlaw spectrum or a blackbody spectrum. In this article we report on observations acquired for both pulsars, namely PSR 1813-36 and \\psr. ",
        "conclusions": "We report the discovery of X-ray emission from a long-period, (1.24s) radio-pulsar with a mild rotational energy loss-rate. Expressed in numbers, the measured power-law fit $L_{X}$=2.25$\\times$10$^{31}$ \\ergsec which should be compared to the 0.1\\% of $\\dot{E}$, a robust formula which gives a good fit to the rotation-powered pulsars (Becker \\& Tr\\\"{u}mper (1997)). The 0.1\\% of $\\dot{E}$ is $10^{29}$\\ergsec. Effectively the observed X-ray Luminosity is 100 times more than that predicted by the magnetospheric model. We should also note that the blackbody fit luminosity is about the same as the powerlaw luminosity. What is the  source of this extra luminosity? \\psr is an extreme pulsar long-period, small $\\dot{E}$. The physical process producing the X-rays may include  reheating due to vortex creep, accretion from a fall-back disc, magnetic decay like the AXPs or accretion from the ISM. Also the luminosity derived from the blackbody fit ($L_{X}$=1.13$\\times$10$^{31}$\\ergsec) and the characteristic age ($P/2\\dot{P}$=2.8$\\times$10$^{6}$ yr) of the pulsar puts it on the standard cooling curve, see Figure 3, although it is on the fast-falling, photon-cooling dominated part.  Deeper observations, to produce sufficient counts with better timing accuracy, are needed in order to resolve the issues stated above. \\cha observations have focused our attention on this extreme pulsar."
    },
    "0212/astro-ph0212221_arXiv.txt": {
        "abstract": "{We study the distribution of bright star-forming complexes in a homogeneous sample of 72 late-type (``irregular'') dwarf galaxies located within the 10 Mpc volume. Star-forming complexes are identified as bright lumps in $B$-band galaxy images and isolated by means of the unsharp-masking method. For the sample as a whole the radial number distribution of bright lumps  largely traces the underlying exponential-disk light profiles, but peaks at a 10 percent smaller scale length. Moreover, the presence of a tail of star forming regions out to at least six optical scale lengths provides evidence against a systematic star formation truncation within that galaxy extension. Considering these findings, we apply a scale length-independent concentration index, taking into account the implied non-uniform random spread of star formation regions throughout the disk. The number profiles frequently manifest a second, minor peak at about two scale lengths. Relying on a two-dimensional stochastic self-propagating star formation model, we show these secondary peaks to be consistent with triggered star formation; for a few of the brighter galaxies a peculiar peak distribution is observed that is conceivably due to the onset of shear provided by differential rotation.  On scales between 100 and 1000 pc, and by taking into account exponential-disk structure, bright lumps reveal cluster dimensions between 1.3 and 2, with a weak trend to higher dimensions for brighter galaxies. Cluster dimension weakly anticorrelates with the lumpiness index (the fraction of the total galaxy light due to the light contributed by the lumps), the latter index showing no dependence on luminosity. Lump spreading within the disk, as measured by the concentration index,  and lump clustering, as given by the cluster dimension, are not linked to each other. Interpreting cluster dimension in terms of porosity of a self-similar intragalactic medium, we derive a relation between current star formation rate, scale length, and porosity. ",
        "introduction": "\\begin{table*}[btp] \\caption[]{Summary of terms} \\label{summary} $$ \\begin{array}{p{0.2\\linewidth}p{0.76\\linewidth}} \\hline\\hline \\noalign{\\smallskip} Lumps: & bright residual features seen in galaxy images after a median filtered version is subtracted from\\\\ &  the original image; synonymous expressions: bright spots or knots;  physical correspondents: star-\\\\ & forming complexes encompassing H$\\;$II regions, young star clusters, and stellar associations;\\\\ Lumpiness index $\\Chi$:   & fractional flux or ratio of the flux due to the lumps within the residual image and the total galaxy\\\\ & light of the original image.\\\\ Cluster dimension $D$:    & correlation dimension for a discrete set of lump centers in a plane, i.e. within a radius $r$ around\\\\ & a typical lump there are $n \\propto r^D$ other lumps; no weighting for lump size or luminosity. If consi-\\\\ & dered as an indirect measure for a three dimensional medium's fractal dimension, $D$ may be\\\\ & related to the volume filling factor of the empty regions, called porosity.\\\\ Concentration index $CI$: & concentration index as the ratio of lump centers in an inner circle and lump centers in an outer\\\\ & annulus, normalized according to some prescription; no weighting for lump size or luminosity.\\\\ \\noalign{\\smallskip} \\hline\\hline \\end{array} $$ \\end{table*} \\begin{table*}[tbp] \\caption[t]{Galaxy and bright-lump data for the 72 irregular dwarf galaxies of our sample} \\scriptsize $$ \\begin{array}{llrcrrrccrcccccc} \\hline \\noalign{\\smallskip} {\\rm galaxy}&{\\rm type} & M_B &\\mu_0^B & R_d & R_{25} &{\\rm b/a} & v_{\\rm rot} & {\\rm Ref.}^\\dagger & N & R_l/R_d & R_{1st}/R_d & R_{2nd}/R_d & CI(R_{25}) & D & \\Chi \\\\ \\noalign{\\smallskip} \\;\\;\\;\\;(1) & (2) & (3)\\;\\; & (4) & (5) & (6) & (7) & (8) & (9) & (10) & (11) & (12) & (13) & (14) & (15) & (16)\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} {\\rm DDO53}         & {\\rm Im}      & -12.84 & 23.06 & 281 & 544 & 0.87 & \\dots & I   &  12 & 1.12 & 1.5 & 3.0 & 0.37 &\\dots& 0.147 \\\\ {\\rm UGC4483}       & {\\rm Im}      & -12.66 & 22.15 & 182 & 578 & 0.50 & 22    & I   &   8 &\\dots & 1.0 & 0.0 & 1.26 &\\dots& 0.101 \\\\ {\\rm Kar54}         & {\\rm Im}      & -14.68 & 23.89 &1283 &1414 & 0.85 & 28    & I   &   1 &\\dots &\\dots&\\dots&\\dots &\\dots& 0.019 \\\\ {\\rm UGC4998}       & {\\rm Im}      & -15.78 & 22.50 & 843 &1879 & 0.50 & \\dots & I   &  10 & 1.14 &\\dots&\\dots& 1.13 &\\dots& 0.058 \\\\ {\\rm HoI}           & {\\rm Im}      & -15.44 & 22.75 &1007 &2089 & 0.83 & 28    & I   &  85 & 0.90 & 2.0 & 1.0 & 0.81 & 1.30& 0.180 \\\\ {\\rm BK1N}          & {\\rm Im}      & -12.70 & 24.49 & 529 &1832 & 0.42 & \\dots & I   &   0 &\\dots &\\dots&\\dots&\\dots &\\dots& 0.035 \\\\ {\\rm NGC2976}       & {\\rm Sd}      & -17.37 & 20.52 & 807 &2988 & 0.42 & 54    & I   & 218 & 0.36 & 2.0 &\\dots& 0.49 & 1.59& 0.077 \\\\ {\\rm UGC5423}       & {\\rm BCD}     & -13.76 & 22.30 & 315 & 694 & 0.67 & 27    & I   &   4 &\\dots &\\dots&\\dots&\\dots &\\dots& 0.051 \\\\ {\\rm DDO82}         & {\\rm Sm/BCD}  & -14.75 & 22.73 & 593 &1209 & 0.56 & \\dots & I   &   5 &\\dots &\\dots&\\dots&\\dots &\\dots& 0.045 \\\\ {\\rm DDO87}         & {\\rm Im}      & -14.18 & 24.13 & 929 & 746 & 1.00 & 66    & I   &   7 &\\dots & 1.5 & 4.5 &\\dots &\\dots& 0.069 \\\\ {\\rm Kar73}         & {\\rm Im}      & -10.81 & 24.53 & 228 & 120 & 0.67 & \\dots & I   &   1 &\\dots &\\dots&\\dots&\\dots &\\dots& 0.030 \\\\ {\\rm UGC8215}       & {\\rm Im}      & -12.71 & 22.29 & 234 & 593 & 0.74 & 14    & III &   0 &\\dots &\\dots&\\dots&\\dots &\\dots& \\dots \\\\ {\\rm DDO167}        & {\\rm Im}      & -12.34 & 22.64 & 248 & 565 & 0.60 & 17    & III &  10 & 0.78 & 1.5 &\\dots& 0.55 &\\dots& 0.083 \\\\ {\\rm DDO168}        & {\\rm Im}      & -14.97 & 21.89 & 715 &2048 & 0.40 & 31    & III &  70 & 0.58 & 1.0 & 2.0 & 1.47 & 1.31& 0.110 \\\\ {\\rm DDO169}        & {\\rm Im}      & -15.30 & 22.15 &1032 &2692 & 0.32 & 26    & III &  15 & 0.74 & 1.0 &\\dots& 1.92 &\\dots& 0.097 \\\\ {\\rm UGC8508}       & {\\rm Im}      & -13.69 & 21.24 & 293 &1005 & 0.56 & 26    & III &  33 & 0.96 & 2.0 &\\dots& 0.63 & 1.82& 0.042 \\\\ {\\rm NGC5229}       & {\\rm Sd}      & -14.44 & 20.51 & 482 &2007 & 0.19 & 56    & III &  30 & 1.10 & 1.0 &\\dots& 1.56 & 1.56& 0.120 \\\\ {\\rm NGC5238}       & {\\rm Sdm}     & -15.03 & 21.90 & 544 &1507 & 0.68 & 19    & III &  22 & 0.46 & 0.5 &\\dots& 6.27 & 1.56& 0.086 \\\\ {\\rm DDO181}        & {\\rm Im}      & -13.30 & 22.00 & 372 &1064 & 0.44 & 21    & III &  44 & 1.11 & 0.5 & 3.5 & 0.99 & 1.67& 0.207 \\\\ {\\rm DDO183}        & {\\rm Im}      & -13.90 & 21.89 & 337 &1016 & 0.75 & 16    & III &   7 &\\dots &\\dots&\\dots&\\dots &\\dots& 0.035 \\\\ {\\rm UGC8833}       & {\\rm Im}      & -11.95 & 22.41 & 177 & 425 & 0.71 & 19    & III &   2 &\\dots &\\dots&\\dots&\\dots &\\dots& 0.020 \\\\ {\\rm HoIV}          & {\\rm Im}      & -15.95 & 22.11 &1634 &4442 & 0.26 & 39    & III &  65 & 1.06 & 1.0 & 4.0 & 0.71 & 1.78& 0.175 \\\\ {\\rm NGC5474}       & {\\rm Scd}     & -17.52 & 20.75 & 992 &3741 & 0.96 & 41    & III & 198 & 1.41 & 1.5 & 4.0 & 0.50 & 1.55& 0.109 \\\\ {\\rm NGC5477}       & {\\rm Sm}      & -15.24 & 21.59 & 559 &1763 & 0.70 & 32    & III &  37 & 1.01 & 2.0 & 1.0 & 0.49 & 1.61& 0.121 \\\\ {\\rm DDO190}        & {\\rm Im}      & -15.17 & 20.99 & 382 &1393 & 0.88 & 44    & III &  24 & 1.02 & 1.0 & 0.0 & 1.23 & 1.49& 0.030 \\\\ {\\rm DDO194}        & {\\rm Im}      & -15.00 & 23.03 & 981 &1814 & 0.64 & 38    & III &   0 &\\dots &\\dots&\\dots&\\dots &\\dots& \\dots \\\\ {\\rm UGC6541}       & {\\rm Sm/BCD}  & -13.40 & 22.49 & 294 & 677 & 0.53 & 17    & IV  &   7 &\\dots & 0.5 &\\dots&\\infty&\\dots& 0.037 \\\\ {\\rm NGC3738}       & {\\rm Irr}     & -15.80 & 22.40 & 685 &1544 & 0.79 & 69    & IV  &  74 & 0.28 & 0.0 &\\dots& 12.3 & 1.95& 0.062 \\\\ {\\rm NGC3741}       & {\\rm Im/BCD}  & -13.34 & 21.71 & 207 & 609 & 0.77 & 38    & IV  &  14 & 1.04 & 1.0 & 2.0 & 1.95 &\\dots& 0.029 \\\\ {\\rm DDO99}         & {\\rm Im}      & -14.51 & 23.09 & 904 &1560 & 0.47 & 23    & IV  &  70 & 0.83 & 0.5 & 2.5 & 1.85 & 1.69& 0.132 \\\\ {\\rm NGC4068}       & {\\rm Sm/BCD}  & -15.69 & 22.89 & 950 &2089 & 0.57 & 30    & IV  & 111 & 0.50 & 1.5 &\\dots& 0.64 & 1.34& 0.243 \\\\ {\\rm NGC4163}       & {\\rm BCD}     & -14.06 & 22.68 & 422 & 597 & 0.66 & 17    & IV  &  17 & 1.14 & 0.5 &\\dots& 1.88 &\\dots& 0.076 \\\\ {\\rm UGC7298}       & {\\rm Im}      & -13.71 & 22.44 & 413 & 959 & 0.67 & 19    & IV  &   0 &\\dots &\\dots&\\dots&\\dots &\\dots& \\dots \\\\ {\\rm NGC4248}       & {\\rm IBm}     & -16.33 & 21.57 &1066 &3313 & 0.43 & 42    & IV  &  16 & 1.02 & 0.5 &\\dots&\\infty&\\dots& 0.072 \\\\ {\\rm DDO127}        & {\\rm Sm}      & -14.39 & 24.72 &1003 &1455 & 0.51 & 32    & IV  &   3 &\\dots &\\dots&\\dots&\\dots &\\dots& 0.039 \\\\ {\\rm UGC7639}       & {\\rm dS0/BCD} & -15.58 & 22.57 &1015 &2184 & 0.54 & 25    & IV  &  64 & 0.54 & 0.5 &\\dots& 4.24 & 1.48& 0.041 \\\\ {\\rm UGC288}        & {\\rm Im}      & -13.82 & 22.98 & 441 & 799 & 0.45 & 25    & VI  &   1 &\\dots &\\dots&\\dots&\\dots &\\dots& 0.000 \\\\ {\\rm UGC685}        & {\\rm Im/BCD}  & -14.92 & 21.96 & 509 &1392 & 0.63 & 38    & VI  &   1 &\\dots &\\dots&\\dots&\\dots &\\dots& 0.019 \\\\ {\\rm UGC1281}       & {\\rm Sd}      & -15.83 & 20.90 &1125 &3725 & 0.16 & 50    & VI  &  94 & 0.70 & 0.5 & 2.0 & 1.96 & 1.69& 0.140 \\\\ {\\rm NGC1156}       & {\\rm IBm}     & -17.68 & 20.31 & 922 &3971 & 0.62 & 55    & VI  &  63 & 0.99 & 1.0 & 2.0 & 0.98 & 1.80& 0.118 \\\\ {\\rm UGC2684}       & {\\rm Im}      & -13.13 & 23.30 & 618 &1100 & 0.38 & 37    & VI  &   1 &\\dots &\\dots&\\dots&\\dots &\\dots& 0.055 \\\\ {\\rm UGC2716}       & {\\rm Sm}      & -15.08 & 23.24 &1047 &1673 & 0.47 & 27    & VI  &   7 &\\dots & 0.0 &\\dots&\\infty&\\dots& 0.044 \\\\ {\\rm UGC2905}       & {\\rm Im}      & -14.41 & 21.30 & 277 & 889 & 0.63 & 26    & VI  &   1 &\\dots &\\dots&\\dots&\\dots &\\dots& 0.016 \\\\ {\\rm UGC3303}       & {\\rm Sd}      & -15.90 & 23.07 &1718 &3031 & 0.61 & 79    & VI  &   7 &\\dots &\\dots&\\dots&\\dots &\\dots& 0.048 \\\\ {\\rm PGC17716}      & {\\rm SBd}     & -16.79 & 21.90 &1310 &3207 & 0.76 & 51    & VI  &   5 &\\dots &\\dots&\\dots& 4.11 &\\dots& 0.068 \\\\ {\\rm A0554+07}      & {\\rm Im}      & -12.25 & 23.46 & 254 & 366 & 0.97 & 27    & VI  &   0 &\\dots &\\dots&\\dots&\\dots &\\dots& \\dots \\\\ {\\rm UGC3476}       & {\\rm Im}      & -14.26 & 21.54 & 484 &1477 & 0.26 & 47    & VI  &   5 &\\dots &\\dots&\\dots& 3.18 &\\dots& 0.082 \\\\ {\\rm UGC3600}       & {\\rm Im}      & -13.53 & 22.88 & 672 &1315 & 0.29 & 39    & VI  &   2 &\\dots &\\dots&\\dots&\\dots &\\dots& 0.049 \\\\ {\\rm NGC2337}       & {\\rm IBm/BCD} & -16.39 & 21.59 & 801 &2637 & 0.62 & 74    & VI  &  33 & 0.91 & 1.0 & 2.0 & 2.15 & 1.66& 0.094 \\\\ {\\rm UGC4115}       & {\\rm Im}      & -13.51 & 22.66 & 493 &1054 & 0.44 & 44    & VI  &  16 & 0.84 & 1.5 & 3.5 & 0.00 &\\dots& 0.077 \\\\ {\\rm NGC2537}       & {\\rm Sm/BCD}  & -16.60 & 22.10 & 711 &1911 & 0.91 & 102   & VI  &  53 & 0.47 & 1.0 &\\dots& 3.34 & 1.59& 0.079 \\\\ {\\rm DDO64}         & {\\rm Im}      & -13.95 & 22.56 & 636 &1532 & 0.39 & 42    & VI  &  17 & 1.83 & 1.5 &\\dots& 0.65 &\\dots& 0.130 \\\\ {\\rm ESO 473-G024}  & {\\rm Im}      & -13.74 & 22.00 & 454 &1179 & 0.46 & 17    & VII &  22 & 0.34 & 1.5 &\\dots& 0.10 & 1.45& 0.196 \\\\ {\\rm ESO 115-G021}  & {\\rm Sm}      & -15.18 & 21.63 &1164 &3455 & 0.15 & 50    & VII & 372 & 0.85 & 1.0 &\\dots& 1.24 & 1.55& 0.257 \\\\ {\\rm ESO 154-G023}  & {\\rm Sm}      & -16.23 & 20.10 &1141 &5635 & 0.22 & 56    & VII & 318 & 1.54 & 2.0 & 3.5 & 0.39 & 1.62& 0.204 \\\\ {\\rm NGC 1311}      & {\\rm Sm}      & -15.69 & 21.74 & 884 &2587 & 0.30 & 42    & VII &  29 & 0.56 & 0.5 &\\dots& 4.02 & 1.78& 0.107 \\\\ {\\rm IC 1959}       & {\\rm Sdm}     & -15.92 & 19.25 & 563 &2485 & 0.24 & 58    & VII & 225 & 2.29 & 2.5 & 0.5 & 0.28 & 1.56& 0.142 \\\\ {\\rm IC 2038}       & {\\rm Sd}      & -14.47 & 21.63 & 596 &1931 & 0.29 & 33    & VII &  33 & 0.41 & 0.5 &\\dots& 4.10 & 1.76& 0.066 \\\\ {\\rm NGC 1800}      & {\\rm Sm/BCD}  & -16.25 & 21.64 & 683 &2086 & 0.57 & 39    & VII &  59 & 0.55 & 0.5 &\\dots& 7.78 & 1.60& 0.083 \\\\ {\\rm AM 0521-343}   & {\\rm Im}      & -14.24 & 21.60 & 358 &1020 & 0.67 & 42    & VII &  11 & 1.14 & 2.0 & 1.0 & 0.49 &\\dots& 0.058 \\\\ {\\rm ESO 555-G028}  & {\\rm Im}      & -13.79 & 23.43 & 870 &1399 & 0.37 & 54    & VII &   7 &\\dots & 1.0 &\\dots& 1.11 &\\dots& 0.065 \\\\ {\\rm ESO 489-G056}  & {\\rm Im}      & -12.30 & 23.00 & 186 & 377 & 0.72 & 19    & VII &  22 & 0.78 & 1.0 &\\dots& 0.82 & 1.47& 0.094 \\\\ {\\rm ESO 490-G017}  & {\\rm Im}      & -14.18 & 21.30 & 269 & 918 & 0.79 & 28    & VII &  36 & 0.90 & 0.5 & 1.5 & 4.21 & 1.47& 0.076 \\\\ {\\rm ESO 308-G022}  & {\\rm Im}      & -13.71 & 23.81 & 609 & 683 & 0.90 & 33    & VII &  11 & 1.26 & 0.5 & 2.0 & 0.46 &\\dots& 0.117 \\\\ {\\rm PGC 20125}     & {\\rm Im}      & -13.46 & 23.75 & 676 & 502 & 0.70 & 36    & VII &  61 & 0.96 & 1.5 & 0.0 & 2.16 & 1.29& 0.292 \\\\ {\\rm ESO 558-PN011} & {\\rm Im}      & -16.30 & 22.08 & 825 &2339 & 0.71 & 68    & VII &  49 & 0.25 & 1.0 &\\dots& 18.7 & 1.38& 0.120 \\\\ {\\rm ESO 059-G001}  & {\\rm Im}      & -14.49 & 22.20 & 536 &1288 & 0.71 & 46    & VII &  60 & 0.70 & 1.0 &\\dots& 1.20 & 1.70& 0.149 \\\\ {\\rm ESO 006-G001}  & {\\rm Im}      & -14.93 & 22.12 & 504 &1271 & 0.85 & \\dots & VII &  22 & 0.41 & 0.0 &\\dots& 9.97 & 1.36& 0.104 \\\\ {\\rm UGCA 148}      & {\\rm Im}      & -14.09 & 21.40 & 342 &1158 & 0.72 & 36    & VII &   6 &\\dots & 1.5 &\\dots& 0.34 &\\dots& 0.038 \\\\ {\\rm UGCA 153}      & {\\rm Sm/BCD}  & -14.21 & 23.32 & 937 &1427 & 0.42 & 52    & VII &  23 & 0.50 & 0.5 &\\dots& 0.77 & 1.44& 0.187 \\\\ {\\rm NGC 2915}      & {\\rm Sm}      & -16.61 & 21.26 & 673 &2290 & 0.56 & 62    & VII &  81 & 0.36 & 0.0 &\\dots& 7.03 & 1.98& 0.091 \\\\ {\\rm UGCA 193}      & {\\rm Sdm}     & -15.15 & 21.98 &1480 &4138 & 0.11 & 55    & VII &  23 & 0.56 & 0.5 &\\dots& 2.36 & 1.24& 0.284 \\\\ \\noalign{\\smallskip} \\hline \\end{array} $$ \\begin{list}{}{}  \\item[$^\\dagger$]References to the published papers of our series on {\\it Structure and stellar content of dwarf galaxies}: $I$ - Bremnes et al. (1998); $III$ - Bremnes et al. (1999); $IV$ - Bremnes et al. (2000); $VI$ - Barazza et al. (2001); $VII$ - Parodi et al. (2002). \\end{list} \\end{table*}  Dwarf irregular galaxies exhibit peculiar morphologies that are dominated by the flashy but seemingly irregular presence of star-forming regions.  Unlike with more massive disk systems that modulate star formation by spiral density waves, dwarf irregulars in the field --- or non-interacting irregulars in general --- constitute ideal testbeds for the study of genuine processes regulating local and global star formation and consequently of galactic evolution. A review addressing several key questions  concerning large-scale star formation in irregular galaxies is given by Hunter (1997), and an evaluation among  simple models for the onset of star formation in irregulars  is provided by Hunter, Elmegreen, \\& Baker (1998).  Important clues as to hidden constraints shaping the heterogeneous appearance of dwarf irregular galaxies may emerge from detailed investigations concerning the spatial distribution of star-forming regions (e.g., Feitzinger \\& Braunsfurth 1984). In recent years several studies on the distribution of H$\\;$II regions and of young, compact star clusters in late-type spirals and in dwarf irregular galaxies have appeared (Telles, Melnick, \\& Terlevich 1997; Brosch, Heller, \\& Almoznino 1998; Elmegreen \\& Salzer 1999; Heller et al. 2000; Roye \\& Hunter 2000; Billett, Hunter, \\& Elmegreen 2002). Applying measures like concentration, asymmetry, and fractional-luminosity indices these authors found the star forming regions to be distributed rather randomly, with some tendency to central concentrations particularly for star-bursting systems.  In this paper  we extend these previous studies by analysing the distribution of bright spots or lumps in the $B$-band images of a sample of 72 late-type (``irregular'') dwarf galaxies. The general equivalence of bright lumps in H$\\alpha$ and in broad band blue images as tracers of star formation complexes can be appreciated by comparing galaxy images filtered at the two corresponding wavelengths (e.g., Elmegreen \\& Salzer 1999; Sparke \\& Gallagher 2000, pp. 139 and 229). With our homogeneous and relatively large sample at hand we aim at comparing three morphological indices with each other (as applied to lumps within irregulars), none of which has been previously reported in the form presented here or within our context. Indices often serve as the quantitative counterparts to qualitative physical concepts. In particular, lump spreading within a galaxy may be described by {\\it concentration indices} of different apertures; lump clustering may be represented by the {\\it correlation dimension} for the two-dimensional lump distribution; finally, the {\\it lumpiness} (or flocculency) of a galaxy may be measured by means of some fractional light index. Table 1 gives a summary of terms and indices that will be more carefully introduced in the subsequent sections.  Comparing relations among morphological indices, we may deepen our insights into the various mechanisms responsible for the morphology of irregulars. The interstellar matter of dwarf irregulars with different global properties may be different (e.g., metallicity, mean gas density, turbulence, gravitational potential), implying differences in the conditions for the formation of stars. Thus aspects concerning the abundance and distribution of star-forming regions within the galaxies, like clustering and star formation rates, may turn out to vary correspondingly. We adress this issue by means of the clustering parameter. Another goal of the paper is to contribute to the discussion concerning the influence of shear due to differential rotation on star formation in gas-rich late-type galaxies (Roye \\& Hunter 2000; Elmegreen, Palou\\v{s}, Ehlerov\\'{a} 2002).  Based on cellular automata simulations we introduce a possible criterium to be checked for in radial lump number distributions; we show that a few of our galaxies indeed meet the criterium, but more research is needed for conclusive results.  The paper is organised as follows. Section 2 gives an overview of the galaxy sample used and presents a table with basic galaxy parameters as well as the parameter values deduced in the subsequent sections. Section 3 describes the adopted lump detection method and introduces a first index, the lumpiness index. Section 4 presents the radial number and  number density distributions for the bright lumps of all the galaxies. A concentration index that is normalized according to the galaxies being exponential-disk systems is introduced and applied to the bright lump distribution. We then extensively discuss the possible reason for a peculiarity seen in the lump number distribution, namely the occurence and relative locations of major and minor peaks. In Section 5 we determine the cluster or correlation dimension for the two-dimensional lump patterns, relate it to galaxy absolute magnitude, and deduce a model mean porosity that is linked to the current  star formation rate. In Section 6 we check for relations among the three indices introduced in the previous sections. We end with a discussion and the conclusions in Section 7. Image processing was performed throughout within the IRAF\\footnote{IRAF is distributed by the National Optical Astronomy Observatories; http://iraf.noao.edu.}  package.  \\begin{figure*}[th] \\begin{center} \\includegraphics[width=85mm]{MS3024f01a.ps} \\includegraphics[width=85mm]{MS3024f01b.ps} \\end{center} \\caption[]{ESO 473-G024. {\\bf Left:} $B$-band image, taken at the 1.5-m Danish Telescope at La Silla, Chile. {\\bf Right:} Residual image after subtracting from the original image its median filtered version.} \\label{eso473} \\end{figure*} ",
        "conclusions": "Regarding the azimuthally integrated, radial distribution of bright lumps --- corresponding to star-forming complexes --- in dwarf irregular galaxies, we find them non-uniformly distributed. While in individual galaxies the number distribution is non-monotonic and rugged, the summed-up distribution for all galaxies of our sample manifests the hidden constraint, which is a $r\\,e^{-r}$-distribution closely tracing the underlying older population.  More precisely, in terms of radial number density distribution the lumps follow an exponential decay with scale length about 10 percent smaller on average than that of the blue continuum light. This is consistent with studies of the radial distribution of H$\\,$II regions in intermediate-type spiral galaxies (Athanassoula et al. 1998). {\\it The fact that each component of the average disk --- from star-forming site number density to surface brightness of the total light --- is approximately exponential is a hint that luminous exponential disks are born rather than made, consistent with the accretion scenario for the viscous evolution of galaxy disks} (e.g., Ferguson \\& Clarke 2001).  Star-forming complexes in irregular dwarf galaxies can be found out to large radii, as already emphasized by Schulte-Ladbeck \\& Hopp (1998), Brosch et al. (1998), and Roye \\& Hunter (2000). {\\it The presence of a tail in the accumulated radial number distribution of star forming regions out to at least six optical scale lengths (Fig. 4)  indicates that the distributions of dwarf irregulars are truncated at rather low gas density thresholds for star formation. (van Zee et al. 1997, Hunter et al. 1998, Pisano et al. 2000); this seems to be different with many spiral galaxies for which sharp to weak truncations, starting at galactocentric radii of 2-4 near-infrared or 3-5 optical scale lengths, have persistently been reported (e.g., recently, Florido et al. 2001, Kregel et al. 2002, Pohlen et al. 2002).}  Beside the presence of main or primary peaks in the radial lump number distribution at slightly less than one optical scale length on average, there is --- contrary to the expectations for exponential-disk systems --- the  frequent occurence of secondary peaks at about two scale lengths. As simple simulations show, it is consistent with the idea of triggered star formation based on a stochastic self-regulation scenario. However, some of the brighter, larger galaxies exhibit pronounced primary peaks at two scale lengths and show  minor, secondary peaks around one scale length. For these galaxies with a reversed peak pattern the simulations indicate that shear-induced star formation around the disk's turnover to differential rotation could be at work; we feel this issue worth a deeper investigation: given a lump statistics based on higher-resolution images and linked to detailed rotational velocity data, and possibly supplemented with information on large-scale magnetic field structures, this may lead to some subtle but decisive insights related to star formation in irregular galaxies. Also, the possible role of bars or bar-like central features should be carefully considered.  However, as none of our four galaxies with principal peaks appering around two scale lengths is classified as ``barred'', and because Roye \\& Hunter (2000) did not see a preferential location of H$\\;$II regions towards the ends of bars in the two candidate galaxies of their sample, we do not consider this mechanism as being effective in shaping the number distribution of lumps.  The observation that the scale lengths are the larger the older the underlying respective population is, goes in hand with the finding that in star-forming dwarf galaxies the oldest populations are also the most extended ones (Gardiner \\& Hatzidimitriou 1992, Minniti \\& Zijlstra 1996, Minniti et al. 1999, Harris \\& Zaritsky 1999). The common interpretation is that of an age-related dispersion of stars. Because dynamical disk heating tends to saturate at a fixed velocity dispersion (Freeman \\& Bland-Hawthorn 2002) the amount of radial spread introduced by dynamical heating  is expected to be larger in smaller galaxies with lower gravitational binding energies (J. Gallagher, private communication). Indeed, the fact noted in Parodi et al. (2002) that with increasing scale length the disk color gradients become systematically less positive and even  start being weakly negative when going from dwarf irregular to low-surface brightness and spiral galaxies ---  which is equivalent to larger galaxies having red-to-blue band scale length ratios below one --- seems to support the idea that the extent of the red component decreases as galaxy mass increases. However, for spirals part of this colour trend is most certainly a metallicity effect.  We provide no investigation of the azimuthal lump distribution; where necessary we assumed axisymmetry.  Applying different asymmetry indices to dwarf and normal irregulars, Heller et al. (2000) and Roye \\& Hunter (2000) obtained opposing results that were dependent on the definition of the chosen index (asymmetry as a question of perception). Also, it was found that while the conveniently applied rotational asymmetry index may be used as a first discriminator between (distant) elliptical and spiral/irregular galaxies (Schade et al. 1995, Conselice et al. 2000), this parameter actually is strongly dependent on recent star formation (Takamiya 1999, Mayya \\& Romano 2001) and thus must be  correlated with the lumpiness index as adopted in the present paper. As we have no data  for the star formation rate of our galaxies, an immediate comparision cannot be offered, however.  Applying a concentration index $CI$ that is normalized according to the exponential-disk structure of a mean lump distribution leads to consistent results for varying aperture sizes. It may also remove the discrepancy found by Heller et al. (2000) between the $CI_{H\\alpha}$ values for actual galaxies and the lower ones for simulated galaxies with random  star formation region positions. Roye \\& Hunter (2000) pointed out an increased scatter of concentration indices for faster rotating galaxies of their sample; we no longer see this effect with our larger sample.  Comparing concentration ($CI$) with lumpiness (\\Chi) we find the galaxies with a high percentage ($>10\\%$) of light stemming from the lumps showing low to moderate concentrations ($CI\\stackrel{_<}{_\\sim}2$), i.e., galaxies with lumps that are widely scatterd within the disk  maintain a higher fraction of the total $B$ luminosity (Fig. \\ref{chi_CI_D}). On the other hand, for very concentrated galaxies ($CI\\stackrel{_>}{_\\sim}2$) less than about 10 percent of the light is due to the lumps; actually, for these galaxies the values for the lumpiness index $\\Chi$ {\\it fluctuate} around the sharp value of 7\\% attributed to most of the (barred) spiral and irregular galaxies observed by Elmegreen \\& Salzer (1999). They suggested the similarity of the blue-band light fraction in complexes for several galaxies of different Hubble types and different total luminosities being due to similar star formation efficiencies.  It remains, however, remarkable that galaxies with very high lump concentrations are {\\rm not} among the galaxies showing high $B$ luminosity fractions of the lumps (Fig. 8, bottom panel).  One may wonder whether some of the nearby, well-known BCDs with central starbursts would agree with this conclusion as well, i.e. whether their central bursts should not dominate the total light content. We therefore examined where the typical starburst dwarf irregular galaxies NGC$\\,$1569 and NGC$\\,$1705 would fill in the $CI$ vs. \\Chi diagram. Based on the 48 brightest central star clusters of NGC$\\,$1569 as compiled in Hunter et al.  (2000) and adopting a distance of 2.5 Mpc and a galaxy absolute magnitude of $M_V=-17.99$ mag, we estimate a mere $\\Chi \\le 0.10$ for NGC$\\,$1569. Similarly, for the super-star cluster dominated amorphous galaxy NGC 1705, the absolute magnitude for the 7 brightest clusters is  about $M_V=-14.1$ according to the data in O'Connell et al. (1994), whereas the galaxy has a magnitude of $M_V=-16.13$ at a distance of 5 Mpc, implying only $\\Chi \\approx 0.15$. Thus independent of the exact concentration indices for the brightest clusters, these two small but representative galaxies with central starbursts would not occupy the empty part of the $CI$ vs. $\\Chi$ diagram. This means that the empty region in this figure is not a selection effect, but may be related to our adopted procedures in determining the corresponding indices.  While the concentration index is a measure for lump spreading, the cluster or correlation dimension provides information on the scaling behaviour for lump-to-lump distances. We found the lump cluster dimensions --- corrected for the effect of radial abundances according to the annulus-integrated exponential distribution --- to lie between 1.3 and 2.0 and to gently correlate with extrapolated central surface brightness and absolute magnitude of the host galaxy. At the same time the cluster dimension is weakly anticorrelated with the lumpiness index.  \\begin{figure*}[t] \\hspace{-5mm} \\includegraphics[width=185mm]{MS3024f09.ps} \\caption{SSPSF simulation of an exponentiated disk with 80 corotating rings and a transition from rigid to strong differential rotation at two scale lengths (ring 32, solid line), after 500 time steps. The three panels show the pattern of newly activated cells (filled circles) together with those created less than 10 time steps ago (dots), the radial number distribution of the active cells (bin width is 8 rings or half a scale length), and the corresponding number density distribution. Note that for the number distribution the peak at one scale length (ring 16) --- the presence of which is expected for non-differentially rotating exponential disks --- is markedly surpassed by a peak around two scale lengths.} \\label{sspsf_galaxy} \\end{figure*}  Cluster dimension (or porosity) as introduced in this paper may be intimately linked to the sizes of the largest, kiloparsec-sized lump compounds (Elmegreen et al. 1996) or to the sizes of star-forming, collapsed expanding shells (e.g., Walter 1999). Elmegreen et al. (1996) found the sizes of the largest compounds within spiral and irregular galaxies to approximately scale with the square root of the galaxy luminosity, or, if normalized by the galaxy semi-major axis $R_{25}$, small galaxies to have slightly larger relative compound diameters than larger galaxies. This was hypothesized to result from gravitational instabilities with the Jeans length or the mean virial density varying with galaxy luminosity. Walter (1999) suggests the holes in dwarf irregulars to be larger than those in late-type spirals because small galaxies have lower masses and correspondingly lower gravitational potentials and lower ambient ISM gas densities, favoring H$\\,$I shells to grow larger. Alternatively, we are temptatively interpreting the cluster dimension of a galaxy in terms of the volume filling factor of empty regions in a fractal medium, i.e. in terms of porosity (defined analogous to Elmegreen 1997), and find the following statistical trends: {\\it (i) fainter galaxies tend to be more porous; (ii) more porous galaxies also have a lumpier  morphology with lower central lump concentrations; (iii) the more porous the galaxy, the lower the star formation rate per kpc$^2$ (equation 4). While these trends are not unexpected, we  provide an objective and  quantitative statistical treatment of these}.  Porosity, or self-similarity, as observed with the bright-lump distribution within dwarf irregular galaxies reflects the self-regulated evolution of the interstellar medium, with stellar feedback and self-gravitation being the main mechanisms. It is thus not to be confused with the still self-similar pattern of dispersed stellar aggregates that initially formed from the fractal interstellar gas, obeying the canonical value of $D$=1.3 for turbulence-driven star formation (Elmegreen \\& Elmegreen 2001). Our sub-galactic areas are showing $D$$\\stackrel{_>}{_\\sim}$1.7, i.e. much higher fractal dimensions. We suspect that in dwarf irregular galaxies either the dispersive redistribution of stars is indeed much more effective than in spiral galaxies (Elmegreen \\& Hunter 2000), and/or that feed-back regulation is responsible for the partially randomized position patterns (e.g., due to the intersection of giant shells).\\\\"
    },
    "0212/astro-ph0212547_arXiv.txt": {
        "abstract": "We present new imaging and spectroscopic observations of the interacting  system Arp 194 ($\\equiv$ UGC 06945  $\\equiv$ VV 126).  The northern component (A194N) is a distorted spiral or ring galaxy likely disrupted by a collision or close encounter with a southern galaxy (A194S). There is evidence that a third galaxy with similar recession velocity is projected on A194N but its role is likely secondary. A194S is connected to A194N by a string of emission knots which motivates our interpretation that the former was the intruder. Three of the knots are easily discernible in B,R, and \\ha\\ images and are assumed to trace the path of the intruder following the encounter, which we estimate occurred a few 10$^8$ yr ago.  Both A194S and N are experiencing strong bursts of star formation: the \\ha\\ luminosity indicates a total star formation rate $\\sim$ 10 \\msol yr$^{-1}$.  The lack of detectable J and K emission from the blobs, along with strong \\ha\\ emission, indicates that an evolved stellar population is not likely to be present. The brightest knot (closest to A194S) shows a star formation rate of $\\approx$1.2 \\msol yr$^{-1}$\\ which, if sustained over a time $\\approx$ 7$\\times 10^7$ yr, could explain the spectral energy distribution. This suggests that the stripped matter was originally predominantly  gaseous. The brightest knot is detected as a FIRST radio source and this is likely the signature  of supernova remnants  related to enhanced star formation.  Motions in the gas between the brightest knot and A194S, traced by an emission line link of increasing radial velocity, suggests infall toward the center of the intruder.  Arp 194 is therefore one of the few galaxies where evidence of ``cross-fueling'' is observed. ",
        "introduction": "Arp 194 (Arp 1966;  $\\equiv$ UGC 06945 $\\equiv$ VV 126; Vorontsov-Vel'Yaminov 1977) is a small-angular-size system of two major components:  a  north-western  galaxy of disrupted morphology (0'.8 x 0'.6;  see Fig. \\ref{fig:ident} and Fig. \\ref{fig:morph}), and an apparently more regular galaxy (0'.35 $\\times$ 0'.35) located $\\approx$ 40 arcsec to the South-East.  From the heliocentric velocity of the Arp 194S ($\\equiv$ A194S hereafter), \\vr $\\approx$ 10500 \\kms\\ (see \\S \\ref{res}), we infer a distance of 175 Mpc for the system (H$_0 \\approx$  60 \\kms\\ Mpc$^{-1}$). Arp 194N  and Arp 194S are therefore separated by a projected linear distance $d_{\\rm p} \\approx$ 34 Kpc (1 arcsec corresponds to $d_{\\rm p}\\approx$ 850 pc). Arp described the Arp 194 system as belonging to the class of ``galaxies with material ejected from nuclei.'' Arp noted further ``outer material connected by thin filament to very hard nucleus.'' Several peculiar features of Arp 194N ($\\equiv$ A194N) have been also discussed in an early study (Metlov 1980).  Theoretical modeling involving a reliable treatment of dissipative phenomena like star formation, as well as  high resolution numerical simulations of stellar and gas motions are still the current frontier in our understanding of interacting galaxies (see e.g., Hearn \\& Lamb 2001; Semelin \\& Combes 2000; Barnes \\& Hernquist 1998; Barnes \\& Hernquist 1996; Mihos \\& Hernquist 1996).  From simulations, we expect a wide range of phenomenological properties due to  interaction. Gravitational forces can marginally enhance the star formation rate (SFR) in an unbound intergalactic encounter as well as give rise to the most luminous bursts of star formation observed in merging systems (see e.g., Krongold et al. 2002, and references therein).  Tidal effects drive noncircular motion in disk galaxies, can strip a significant amount of stars and gas, and even lead to the formation of the so-called tidal-dwarf galaxies (Barnes \\& Hernquist 1992). A related issue is the role of gravitational interaction in fueling nonthermal nuclear activity. Both simulations and observations of single interacting systems are a necessary complement to statistical studies on the frequency of interacting systems among active galaxies, since the latter provide valuable but only circumstantial evidence in favor of interaction as a major driver of nuclear activity (Krongold et al. 2002; see however Schmitt 2001, and references therein).  Collisional ring galaxies are excellent laboratories for studying galactic evolution, global star formation, and the occurrence of gas cross-fueling among galaxies.  The prototype of this class of objects is the Cartwheel galaxy A0035-33. Only few systems similar to this object are known: among the best studied we recall VII~Zw~466, Arp 10, II Zw 28, II Hz 4, Abell 76, LT 36, NGC 985, LT 41 (Appleton \\& Marston 1997).  Other likely cases  are Arp 119 (Hearn \\& Lamb, 2001), Arp 118 (Charmandaris et al. 2001), Arp 284 (Smith \\& Wallin 1992, Smith et al. 1997), and Arp 143 (Appleton et al. 1992). The dominant ring morphology is thought to result from a head-on collision between two galaxies, one of which  traveled close to the spin axis of the other, striking the disk close to its center.  The resulting gravitational perturbation is believed to drive a set of symmetrical waves or caustics through the stellar disk (Lynds \\& Toomre 1976, Theys \\& Spiegel 1976, Toomre 1978, Appleton \\& Struck-Marcell 1987, Struck-Marcell 1990). Computer simulations of such head-on encounters, in which at least one galaxy has a significant gaseous component, show indeed the production of density enhancements and shock waves in the interstellar medium. These high density regions coincide with the location of recent, large scale star formation observed in  ring galaxies.  Fig.  \\ref{fig:ident} and Fig. \\ref{fig:morph} suggest that A194N is a collisionally  induced ring galaxy connected to the past intruder A194S by a string of relatively bright knots (\\S  \\ref{res}).  Several photometric properties indicate star formation rates typical of Starburst galaxies (\\S \\ref{sform}). The surprisingly strong radio emission from A194S and knot A is likely due mainly to supernova remnants which imply even higher star formation rates (\\S \\ref{rsn}) and substantial internal obscuration. Relatively rare systems like Arp 194 where gas motions are induced by a head-on encounter enable us to unambiguously infer the geometry of the system. In this case that information provides evidence for ``cross-fueling'' (see \\S \\ref{int}).  Our results for Arp 194 have intriguing implications  on the analysis and interpretation of galactic superwinds (\\S \\ref{disc}). ",
        "conclusions": "We have analyzed in detail the photometric properties and the kinematics of the strongly interacting system Arp 194. The main constituents of this system are a collisional ring system (A194N) and an intruder (A194S) that have experienced an interpenetrating, head-on encounter a few $ 10^8$ yr ago. We have shown that tidally stripped gas is falling toward the center of the intruder, A194S, and that it is fueling a strong nuclear and circumnuclear Starburst. Arp 194 is therefore one of the few known objects for which convincing evidence of cross-fueling exists. Considering that gas is usually confined in a ``dynamically cold'' configuration in disk galaxies, the Arp 194 case indicates transfer of orbital energy to the internal motion of gas during the the encounter. We suggest that gas motions in superwind galaxies -- which are mostly interacting systems -- could also be affected by the same mechanism.  Several aspects of the Arp 194 system deserve further scrutiny. The morphology of A194N is fairly complex, and the distorted morphology of the ring as well as the bright arc of A194N indicate the possibility of perturbations by a third party. A full coverage with slit spectroscopy will allow to confirm the presence of a second intruder galaxy located approximately in correspondence of the northern side of the ring. High spatial resolution spectroscopy may reveal more complex motions in proximity of the nucleus of A194S. In the interacting galaxy pair NGC 1409/10, the ongoing mass transfer seems to follow a spiraling pattern (Keel 2000; similar considerations may apply to NGC 7592: Hattori et al. 2002). There is no indication of an hidden AGN presence from our data, but this result depends on the resolution of the FIRST survey data. It is possible that higher resolution data may reveal  a high surface brightness compact core. In addition, further radio observations could detect radio supernova events: if $ dn/dt \\approx 0.3$, and if a radio supernova event remains detectable for $\\approx$3 yr, we expect to reveal 1 event at any observing time."
    },
    "0212/astro-ph0212292_arXiv.txt": {
        "abstract": "In the 60s and 70s super-massive central objects (from now onwards SMOs) were thought to be the main source of active galactic nuclei (AGNs) characteristics (luminosities of $L \\approx 10^{12} L_{\\odot}$). The release of gravitational binding energy by the accretion of material on to an SMO in the range of $10^7 - 10^9 M_{\\odot}$ has been suggested to be the primary powerhouse (Lynden-Bell 1969). That rather exotic idea in early time has become common sense nowadays. Not only our own galaxy harbours a few million-solar mass black hole (Genzel 2001) but also many of other non-active galaxies show kinematic and gas-dynamic evidence of these objects (Magorrian et al. 1998) The concept of central super-massive stars (SMSs henceforth) (${\\cal M} \\ge 5 \\times 10^4 M_{\\odot}$, where ${\\cal M}$ is the mass of the SMS) embedded in dense stellar systems was suggested as a possible explanation for high- energy emissions phenomena occurring in AGNs and quasars (Vilkoviski 1976, Hara 1978), such as X-ray emissions (Bahcall and Ostriker, 1975). SMSs and super-massive black holes (SMBHs) are two possibilities to explain the nature of SMOs, and SMSs may be an intermediate step towards the formation of SMBHs (Rees 1984). In this paper we give the equations that describe the dynamics of such a dense star-gas system which are the basis for the code that will be used in a prochain future to simulate this scenario. We also briefly draw the mathematical fundamentals of the code. ",
        "introduction": "To go from stationary to dynamical models we use a gaseous model of star clusters in its anisotropic version. It is based on the basic assumptions that: the system can be described by a one particle distribution function, the secular evolution is dominated by the cumulative effect of small angle deflections with small impact parameters (Fokker-Planck approximation, good for large $N$-particle systems) and that the effect of the two-body relaxation can be modelled by a local heat flux equation with an appropriately tailored conductivity.  \\noindent The first assumption justifies a kinetic equation of the Boltzmann type with the inclusion of a collisional term of the Fokker-Planck (FP) type:  \\begin{equation} \\frac{\\partial f}{\\partial t}+v_{r}\\frac{\\partial f}{\\partial r}+\\dot{v_{r}} \\frac{\\partial f}{\\partial v_{r}}+\\dot{v_{\\theta}}\\frac{\\partial f}{\\partial v_{\\theta}}+\\dot{v_{\\varphi}}\\frac{\\partial f}{\\partial v_{\\varphi}}=\\Bigg ( \\frac{\\delta f}{\\delta t} \\Bigg)_{FP} \\end{equation} \\noindent In spherical symmetry polar coordinates $r$, $\\theta$, $\\phi$ are used and $t$ denotes the time. The vector ${\\bf v} = (v_i), i=r,\\theta,\\phi$ denotes the velocity in a local Cartesian coordinate system at the spatial point $r,\\theta,\\phi$. The distribution function $f$ is a function of, $r$, $t$, $v_r$, $v_\\theta^2+v_\\phi^2$ only due to spherical symmetry. By multiplication of the Fokker-Planck equation (1.1) with various powers of the velocity components we get up to second order a set of moment equations which is equivalent to gas-dynamical equations coupled with Poisson's equation: A mass equation, a continuity equation, an Euler equation (force), radial and tangential energy equations. The system of equations is closed by a phenomenological heat flux equation for the flux of radial and tangential r.m.s. kinetic energy, both in radial direction. ",
        "conclusions": ""
    },
    "0212/astro-ph0212401_arXiv.txt": {
        "abstract": "We have developed a new galactic chemo-dynamical evolution code, called {\\tt GCD+}, for studies of galaxy formation and evolution. This code is based on our original three-dimensional tree N-body/smoothed particle hydrodynamics code which includes self-gravity, hydrodynamics, radiative cooling, star formation, supernova feedback, and metal enrichment. {\\tt GCD+} includes a new Type II (SNe II) and Ia (SNe Ia) supernovae model taking into account the lifetime of progenitor stars, and chemical enrichment from intermediate mass stars. We apply {\\tt GCD+} to simulations of elliptical galaxy formation, and examine the colour-magnitude relation (CMR), the Kormendy relation, and the [Mg/Fe]--magnitude relation of simulation end-products. {\\tt GCD+} is a useful and unique tool which enables us to compare simulation results with the observational data directly and quantitatively. Our simulation confirm the results of Kawata (2001) who uses a simpler chemo-dynamical evolution code. We newly find that radiative cooling becomes more efficient and thus the gas infall rate increases, with decreasing mass of galaxies, which contributes to the slope of the CMR. In addition, the sophisticated treatments of both SNe II and SNe Ia in {\\tt GCD+} show that feedback from SNe Ia plays a crucial role in the evolution of elliptical galaxies. We conclude that the feedback effect of SNe Ia should not be ignored in studying the evolution of elliptical galaxies. ",
        "introduction": "\\label{sintro}  The abundances of elements heavier than helium (``metals'') in interstellar gas and in stars vary systematically from place to place within galaxies. Since heavy elements are believed to have been synthesized in stars, the metal abundances in present-day galaxies should offer a record through which we may trace the history of star formation within galaxies. Since the pioneering work of \\citet{bt72}, studies of chemical evolution have succeeded in connecting observed metal abundances to the history of star formation within galaxies \\citep[e.g.][]{ay87, mt87, tww95, bg97, cmg97, tcbmp98}. Moreover, recently developed semi-analytic models \\citep{wr87, wf91} including chemical evolution make it possible to relate the observed chemical properties of galaxies to cosmological theory \\citep[e.g.][]{kc98, clbf00, on01, spf01, bbfsf02}. These techniques require low computational costs and allow large coverage of parameter space. Although the pure chemical evolution models and the semi-analytic models seem to succeed in explaining the observational properties of galaxies at various redshifts, they inevitably assume a phenomenological model involving a number of parameters to describe the processes of gas accretion rates, radiative cooling, star formation, and supernova feedback within a galactic halo. In addition, structures of the end-products can not be discussed, because there is no information about the dynamical history.  On the other hand, recent advances in computer technology and numerical methods have made it possible to calculate the dynamical and chemical evolution of galaxies self-consistently \\citep[e.g.][]{clc98,mtlc01,cc02}. Three dimensional chemo-dynamical evolution codes have been applied to the following studies over the past five years: \\citet{sm95} studied the chemo-dynamical evolution of disk galaxies and succeeded in distinguishing the chemical properties between halo, bulge, and disk stars \\citep[see also][]{bc00,bc01}. \\citet{rvn96} and \\citet{pb99} took into account metal enrichment by Type Ia supernovae (SNe Ia) as well as Type II supernovae (SNe II), and reproduced the correlation between [O/Fe] and [Fe/H] for stars in the solar neighbourhood. Such self-consistent calculation of chemical evolution allows one to analyse the photometric properties of simulation end-products in combination with stellar population synthesis, and to obtain the luminosity of the end-products without the assumption of an arbitrary mass to light ratio. Taking advantage of such technique, \\citet{sn99}, \\citet{ns00}, and \\citet{ksw00} discussed the zero-point of the Tully-Fisher relation of disk galaxies. \\citet{bs98,bs99,bs00,bs01} studied the dynamical and photometric properties of elliptical galaxies formed by a merger of two disk galaxies. \\citet{mytn97} showed that the galactic wind caused by supernovae (SNe) feedback leads to a positive colour gradient observed in some dwarf elliptical galaxies. Using cosmological simulation, \\citet{ng98}, \\citet{co99}, \\citet{tlmc01}, and \\citet{tvk02} studied metal enrichment of the intergalactic medium \\citep[see also][as a work coupling semi-analytic model to hydro simulations]{ahs01}, and \\citet{nfco01} examined the luminosity function and colour distribution function of galaxies at different redshifts. Since radiative cooling depends on the metallicity, dynamical evolution is quite sensitive to the chemical evolution \\citep{kh98,kpj00}. Thus, we must calculate chemical and dynamical evolution self-consistently.  Using a chemo-dynamical evolution code, \\citet[][hereafter K01b]{dk01b} were able to reproduce the colour-magnitude relation (CMR) of observed elliptical galaxies. K01b showed that SNe feedback affects the evolution of lower mass systems more strongly, and induces a stronger galactic wind in lower mass systems. Such low mass systems have lower metallicities and bluer colours than higher mass systems, as predicted by analytic models \\citep[e.g.][]{ay87,bg97} and more phenomenological numerical simulations \\citep[e.g.][]{rl74,rc84}. Moreover, K01b found that galactic winds are triggered mainly by SNe Ia rather than SNe II, although K01b ignored the lifetime of progenitors of SNe II and SNe Ia (i.e.\\ K01b assumes the instantaneous recycling approximation for SNe II, and all SNe Ia occur simultaneously 1.5 Gyr after stars were born).  We have developed a new chemo-dynamical evolution code, called {\\tt GCD+}, which adopts a more sophisticated chemical evolution model than that discussed in K01b. We have relaxed the instantaneous recycling approximation for SNe II, and adopt a more sophisticated SNe Ia model \\citep[][hereafter KTN00]{ktn00}. Our code takes into account the metallicity dependent lifetime of stars \\citep{tk97}, and metal enrichment from intermediate mass stars \\citep{vdhg97}. The purpose of this paper is to introduce details of {\\tt GCD+}. In addition, as an application of {\\tt GCD+}, we carry out simulations of elliptical galaxy formation similar to the ones of K01b. Then we re-examine the CMR, the Kormendy relation, and the [Mg/Fe]--magnitude relation, which were studied in K01b. Since these observed relations are well-established, and are expected to reflect the chemo-dynamical evolution of elliptical galaxies, this is an important application for {\\tt GCD+}. Next, taking advantage of the sophisticated SNe II and SNe Ia models in {\\tt GCD+}, we clarify the role of SNe Ia, compared to SNe II, in the evolution of elliptical galaxies.  The outline of this paper is as follows. Details of {\\tt GCD+} are described in Section \\ref{scode}. In Section \\ref{sapegf}, we use {\\tt GCD+} to simulate elliptical galaxy formation. We show the galaxy formation model used in this paper briefly in Section \\ref{smodel}. In Sections \\ref{scmr}--\\ref{smgfe}, results of numerical simulations are presented, and are compared to the observed scaling relations. Discussion about the study of elliptical galaxy formation is given in Section \\ref{sdisc}. Finally, we present our conclusions from this paper in Section \\ref{sconc}.  \\begin{figure} \\leavevmode \\epsfxsize=80mm \\epsfbox{cr.ps} \\caption{ Cooling rates as a function of temperature. Each curve corresponds to a cooling rate with different metallicities as indicated. } \\label{cr-fig} \\end{figure} ",
        "conclusions": "\\label{sconc}  We have developed a new chemo-dynamical evolution code called {\\tt GCD+}. The code includes both SNe II and SNe Ia modeling taking into account the lifetime of progenitor stars, and the chemical enrichment from intermediate mass stars. In this paper we have applied the code to simulations of elliptical galaxy formation described as the evolution of an isolated seed galaxy. We have shown that {\\tt GCD+} is a useful and unique tool which enables us to compare simulation results with the observational data directly and quantitatively, and gives us new insights into the effects of dynamics on the chemical evolution of galaxies. Encouraged by the success of this study, we plan to apply {\\tt GCD+} to high-resolution cosmological simulations in the near future. In the real universe, even isolated galaxies at $z=0$ are expected to experience interactions with the another galaxies in the past. In fact, \\citet{ng98} suggests that interactions between galaxies are a dominant mechanism to transport metals from the galaxies to intergalactic medium rather than SNe feedback. Therefore, using simulations covering a cosmological scale, it will be valuable to examine what physical process, or combination of processes, induces the tight observed CMR. In addition, cosmological simulations using {\\tt GCD+} can provide detailed metallicity distributions not only in the stellar component of individual galaxies, but also in the intergalactic medium, such as Lyman-$\\alpha$ clouds and the hot gas in clusters of galaxies. As mentioned in Section \\ref{sfd}, the different chemical elements have progenitors with different masses, and different mass stars have different lifetimes. It is highly probable that the abundance ratios are fairly sensitive to the metal transport mechanism and its epoch. New observational facilities will offer a wealth of information about the distribution of metallicity, as well as the abundance ratio of each element, such as [N/O], [Mg/Fe] and [Si/Fe], with unprecedented accuracy. Studies comparing new observations with our future cosmological simulation results will show clearly which metal transport mechanism is dominant, and thus contribute to understanding the formation and evolution of galaxies as well as clusters of galaxies from a new perspective."
    },
    "0212/astro-ph0212346_arXiv.txt": {
        "abstract": "The interplay among the possible variation of gauge coupling and the inflationary dynamics is investigated in a simplified toy model. Depending upon various parameters (scalar mass, curvature scale at the end of inflation and at the onset of the radiation epoch), the two-point function of the magnetic inhomogeneities grows during the de Sitter stage of expansion and, consequently, large scale magnetic fields are generated. The requirements coming from inflationary magnetogenesis are examined together with the theoretical constraints stemming from the explicit model of the evolution of the gauge coupling. ",
        "introduction": "Suppose  that during a de Sitter stage of expansion the coupling constant of an Abelian gauge field evolves in time. Thus the kinetic term of the gauge field can be written (in four space-time dimensions) as \\begin{equation} S_{\\rm em} = -\\frac{1}{4}\\int d^4 \\, x \\sqrt{- G} f(\\phi) F_{\\alpha\\beta} F^{\\alpha\\beta}, \\label{ac1} \\end{equation} where $G$ is the determinant of the space-time metric, $\\phi$ is a scalar field (which can depend upon space and time) and $g(\\phi)= f(\\phi)^{-1/2}$ is the coupling \\footnote{The Heaviside electromagnetic system of units will be used throughout the investigation. The effective ``electron'' charge will then be given, in the present context, by $ e(\\phi) = e_1 f(\\phi)^{-1/2}$.}. This type of vertex is typical of scalar-tensor theories of gravity \\cite{bg} and of the low energy string effective action \\cite{mmv}. Early suggestions that the Abelian gauge coupling may change over cosmological times were originally made by Dirac \\cite{dir1} (and subsequently discussed in \\cite{dir2} and in \\cite{dir3}) mainly in the framework of ordinary electromagnetism.  The dynamics of the field $\\phi$ will be described by the action of a minimally coupled (massive) scalar. If the field is displaced from the minimum of its potential during inflation, there will be a phase where the field relaxes. Provided the scalar mass is much smaller than the curvature scale during inflation such a phase could be rather long. During the de Sitter stage, the specific form of the expanding background will dictate, through the equations of motion, the rate of suppression of the amplitude of $\\phi$.  While the field relaxes toward the minimum of its potential, energy is pumped from the  homogeneous mode of $\\phi$ to the gauge field fluctuations. The function $f(\\phi)$ can be either an increasing function of $\\phi$ (leading to a decreasing coupling) or a decreasing function of $\\phi$ (leading to an increasing coupling). In both cases, depending upon the parameters of the model,  the two-point correlation function of magnetic inhomogeneities increases during the inflationary stage. This implies that large scale magnetic fields can be potentially generated.  The implications of large scale magnetic fields for cosmology, astrophysics and high-energy physics have been recently reviewed in \\cite{ch}  where various useful references are reported. The interested student will find in \\cite{ch} a morte physical introduction to the subject of large scale magnetic fields. The toy model discussed in the present paper is only an useful exercise suggesting a possible connection between the variation of gauge couplings and the production of large scale magnetic fields.  Recently \\cite{webb} the variation in the fine structure constant has been suggested by observations of absorption lines at different frequencies. The results of these claims can be summarized by saying that $\\alpha_{\\rm em}$ was smaller in the past. Defining, indeed $\\Delta \\alpha_{\\rm em} = \\alpha_{\\rm today} - \\alpha_{\\rm past}$, measurements suggest that \\begin{equation} \\frac{\\Delta \\alpha_{\\rm em} }{\\alpha_{\\rm em} } = (0.72 \\pm 0.18) 10^{-5}. \\end{equation} To avoid confusions, in the present scenario, gauge couplings are {\\em constant} today and their variation occurs prior to big-bang nucleosynthesis (BBN).  The gauge coupling has to be (almost) constant by the time when the Universe is approximately old of one second, namely by the time of BBN. The abundances of light elements are very sensitive to any departure from the standard cosmological model. Hence, fluctuations in the baryon to photon ratio, matter--antimatter domains, anisotropies in the expansion of the four space-time dimensions, can all be successfully constrained by demanding that the abundances of the light elements are correctly reproduced. Following the same logic the variation in the gauge couplings can also be constrained from BBN \\cite{coup}.  The plan of the present contribution is then the following. In Section II the basic ideas concerning the model of evolution of the gauge coupling will be introduced. In Section III bounds coming both from the homogeneous and from the inhomogeneous evolution of $\\vf$ will be described. In Section IV the evolution of the magnetic inhomogeneities will be addressed along the  various stages of the model with particular attention to the role of the two-point function. In Section V the large scale magnetic fields produced in the scenario will be estimated. Section VI contains some concluding remarks. The considerations reported in the present paper extend and complement the results of \\cite{mg}.  Finally, in concluding this Introduction, I wish to remind a couple of useful references of two other lecturers of the Chalonge school. In Ref. \\cite{raf}  an interesting discussion on the possible ambiguities arising in the observations of large scale fields in clusters is reported. In \\cite{danhector} another idea on the generation of magnetic fields is reported.   \\renewcommand{\\theequation}{2.\\arabic{equation}} \\setcounter{equation}{0} ",
        "conclusions": "There are no reasons why the gauge couplings should be constant throughout all the history of the Universe. If they are allowed to change prior to the formation of the light elements they can lead to computable differences in the cosmological evolution.  In the present paper the interplay between inflationary magnetogenesis and the evolution of the Abelian gauge coupling has been addressed. In a phenomenologically reasonable model of inflationary and post-inflationary evolution the relaxation of the gauge coupling leads to a growth in the correlation function of magnetic inhomogeneities. Large scale magnetic fields are then generated. The evolution of the gauge coupling is driven, in the present context, by a massive scalar.  Since the gauge coupling evolves significant changes in the evolution of the plasma can be envisaged. In the present investigation the ordinary MHD equations have been generalized to the case of time evolving gauge coupling but other effects could be envisaged. In particular, the generlization of the full kinetic approach to the case of time evolving ``electron'' charge would be of related interest.  The value of large scale magnetic fields produced with this mechanism has been estimated. For a broad range of parameters the obtained values of the magnetic fields are much larger than the dynamo requirements.  In the present investigation the main assumption has been that the only gauge coupling free to evolve is the Abelian one. Furthermore, the parameters of the model have been chosen in such a way that the gauge coupling is not dynamical by the onset of the electroweak phase transition. It would be interesting to relax both assumptions since they may lead to potential differences with the standard scenarios.  Therefore, in many respects, the present investigation is not conclusive. At the same time it shows that acceptable  models  for the evolution of the gauge couplings can be obtained in a standard cosmological framework."
    },
    "0212/astro-ph0212493_arXiv.txt": {
        "abstract": "Subdwarf B (sdB) stars form the blue end of the horizonal branch. Their peculiar atmospheric abundance patterns are due to diffusion processes. However, diffusion models fail to explain these anomalies quantitatively. From a NLTE model atmosphere analysis of 40 sdB stars, we found that the more luminous (i.e. more evolved) stars have anomalous H$\\alpha$ and HeI\\,6678\\AA\\ line profiles, i.e, the lines are too broad and shallow and may even show some emission. We interpret these anomalies as the signatures of a stellar wind, the first such detection in this class of star (if confirmed).  Mass loss may also explain the peculiar abundance patterns seen in sdB stars. High-quality UV spectra are needed to confirm that these stars do have stellar winds. ",
        "introduction": "Subdwarf B (sdB) stars are core helium burning stars with very thin hydrogen envelopes which all have masses of 0.46-0.50 solar masses and can be identified with models of extreme horizontal Branch (EHB) stars (Heber, 1986).  \\begin{figure} \\vspace{8cm} \\special{psfile=sdb_hrd_ob.ps hscale=40 vscale=40 hoffset=+40 voffset=-60 angle=0} \\caption{(T$_{\\rm eff}$,log\\,g) diagram for the Maxted et al. (2001) sdB sample. Note that most stars are well matched by models for the extreme horizontal branch. Five stars (labelled by their first 4 digits of their PG catalog names) have already evolved from the EHB.}\\label{hrd} \\end{figure}  The atmospheres of sdB stars are typically deficient in helium by more than one order of magnitude, but there is a wide spread from solar He/H to stars with He/H$<$10$^{-4}$. These He abundances are still much too large to be accounted for by diffusion, i.e. by the balance between radiative levitation and gravity (Michaud et al., 1989) which predicts He abundances {\\bf lower} by two orders of magnitudes than the average observed He abundance of He/H$\\approx$10$^{-2}$ (see Fontaine \\& Chayer, 1997). Since the diffusion timescale is small (10$^4$ yrs) compared to the EHB lifetime (10$^8$ yrs), {\\bf no} He should be visible in the radiative atmospheres: the puzzle here is not so much a lack of He but an excess in the sdB atmospheres.  Modelling of metal diffusion in sdB stars is still in its infancy. First attempts have been published by Charpinet et al. (1997) for Fe and Unglaub \\& Bues (2001) for C, N and O but met the observational constraints with very little success (except for iron). ",
        "conclusions": ""
    },
    "0212/astro-ph0212170_arXiv.txt": {
        "abstract": "We present the first strong detection of very high energy $\\gamma$-rays from the close ($z=0.048$) X-ray selected BL Lacertae object 1ES1959+650. Observations were made with the Whipple 10m telescope on Mt. Hopkins, Arizona, using the atmospheric Cherenkov imaging technique. The flux between May and July 2002 was highly variable, with a mean of $0.64\\pm0.03$ times the steady flux from the Crab Nebula and reaching a maximum of five Crab, with variability on timescales as short as seven hours. ",
        "introduction": "Blazars, consisting of flat spectrum radio quasars (FSRQs) and BL Lacertae objects (BL Lacs), are a relatively rare type of active galaxy characterized by extremely variable, non-thermal spectral energy distributions (SEDs). The emission extends from radio to $\\gamma$-ray frequencies and is believed to be produced by a highly relativistic plasma jet closely orientated with the line of sight to the galaxy \\citep{Urry95, Blandford78}. In a $\\nu F_{\\nu}$ representation, the SEDs display two peaks. The lower energy peak is usually attributed to synchrotron radiation from relativistic electrons, while the higher energy peak is thought to be due to inverse Compton scattering by these same electrons. The photons involved in the scattering process may be either synchrotron photons produced in the jet, in the synchrotron self-Compton (SSC) model, or some external population, such as photons emitted from an accretion disc \\citep{Dermer92} or reflected from emission line clouds \\citep{Sikora94}. Alternative explanations for the high energy component invoke pair cascades triggered via pion and pair photoproduction from high energy protons in the jet \\citep{Mannheim93, Mannheim98}, or proton synchrotron radiation \\citep{Aharonian00, Mucke01}.  The EGRET detector on board the CGRO detected in excess of 66 blazars \\citep{Hartman99} above $100\\U{MeV}$; however, only a few have been detected at higher energies ($>~300\\U{GeV}$) by ground based atmospheric Cherenkov telescopes. Markarian~421 \\citep{Punch92} and Markarian~501 \\citep{Quinn96} were the first extragalactic TeV sources to be detected and the behaviour of these two nearby ($z=0.031$ and $z=0.034$ respectively) objects has been closely monitored and studied. Flux variability over two orders of magnitude has been observed, with doubling timescales as short as $\\sim15\\U{minutes}$ \\citep{Gaidos96}, placing strong constraints on the size of the emission region. The $\\gamma$-ray flux is correlated with the X-ray flux \\citep{Buckley96, Catanese97a, Krawczynski97} and, in the case of Markarian 421, with the spectral power law index in the $\\gamma$-ray region \\citep{Krennrich02,Aharonian02a}. Unconfirmed detections of 1ES 2344+514 \\citep{Catanese98} and PKS 2155-304 \\citep{Chadwick99} have also been reported and, most recently, the detection of very weak emission from H1426+428 has been confirmed by three groups \\citep{Horan02, Aharonian02b, Djannati02}. This last object is of particular interest since it is relatively distant ($z=0.129$) and hence the observed flux is expected to be strongly attenuated by absorption via pair production on the extragalactic infra-red background light (EBL) \\citep{Gould67, Stecker92}.  All of the TeV blazars can be classified as high frequency peak BL Lacs (HBLs) \\citep{Padovani95} on the basis of the location of the lower energy peak in their SEDs, and predictions of TeV emission from other nearby sources of this type have driven observations by atmospheric Cherenkov detectors. 1ES1959+650 ($z=0.048$) was first suggested as a good TeV candidate of this type by \\citet{Stecker96}, who used simple scaling arguments to predict it as the third strongest TeV blazar, after Markarian~421 and Markarian~501, with a flux prediction of $1.9\\times10^{-11}\\UU{cm}{-2}\\UU{s}{-1}$ above $300\\U{GeV}$. More recently \\citet{Costamante02} have predicted fluxes above $300\\U{GeV}$ of $7.5\\times10^{-11}\\UU{cm}{-2}\\UU{s}{-1}$, based on a simple phenomenological parameterization of the SED adapted from \\citet{Fossati98}, and $0.03\\times10^{-11}\\UU{cm}{-2}\\UU{s}{-1}$ using a homogeneous, one-zone synchrotron self-Compton model \\citep{Ghisellini02}.  In fact, 1ES1959+650 is not a particularly extreme example of an HBL. Measurements with BeppoSAX in 1997 \\citep{Beckmann02} placed the lower energy SED peak at $10^{15}\\U{Hz}$ ($4\\U{eV}$), as compared to typically $\\sim1\\U{keV}$ for Markarian~421 \\citep{Maraschi99} and $\\sim100\\U{keV}$ for Markarian~501 during its 1997 flare state \\citep{Pian98}, although for both of these sources the position of the peak is known to vary widely depending upon the flux level. \\citet{Giebels02} report on X-ray observations of 1ES1959+650 obtained with the USA(1-16\\U{keV}) and RXTE (2-16\\U{keV}) missions during 2000, showing threefold increases in the X-ray flux on a timescale of a few days. The flux increase is correlated with spectral hardening, indicative of a shift of the lower energy peak of the SED towards higher frequencies. The source is also unusual among HBLs for its strong rapid optical variability. Observations by \\citet{Villata00} showed rapid flickering, including a decrease of 0.28~magnitudes in four days.  Prior to May 2002, only tentative evidence had been presented for TeV emission from 1ES1959+650. A detection with a statistical significance of $3.9\\U{\\sigma}$ was reported by \\citet{Nishiyama99} based on $57\\U{hours}$ of observations with the Utah Seven Telescope Array. More recently, \\citet{Konopelko02} reported a preliminary detection at $\\sim5\\U{\\sigma}$ for the HEGRA Cherenkov telescope system. Previous observations of 1ES1959+650 with the Whipple 10m telescope produced an upper limit at a flux level of $1.3\\times10^{-11}\\UU{cm}{-2}\\UU{s}{-1}$ above $350\\U{GeV}$ \\citep{Catanese97b}.  Observations of 1ES1959+650 during May-July 2002 with the Whipple 10m telescope resulted in a clear detection of TeV $\\gamma$-ray emission from this object \\citep{IAUC} which was rapidly confirmed by the HEGRA experiment \\citep{IAUC2}. We report here the results of the Whipple 10m observations. ",
        "conclusions": "The detection of TeV $\\gamma$-ray emission from 1ES1959+650 adds another member to the class of TeV blazars, all of which are close BL Lac objects having a low bolometric luminosity and the peaks in their SEDs at high frequency. The rapid flux variability, and the fact that 1ES1959+650 has not been detected during previous observations, indicates that the source was in an unusual flaring state during these observations. The flux level was at times orders of magnitude above the most recent model predictions. Throughout the period of the Whipple observations measurements in the $2-12\\U{keV}$ region by the All-Sky Monitor (ASM) on board RXTE have shown 1ES1959+650 to be active and variable, with daily average fluxes reaching $\\sim20\\U{mCrab}$ in May and July. Target of opportunity observations with the RXTE small field of view instruments were triggered following the $\\gamma$-ray detection and will be reported on elsewhere (Krawczynski, private communication). The rapid flux variability observed at TeV energies implies a small emission region in the jet with a high Doppler factor \\citep{Mattox93,Madejski96,Buckley96}. The contemporaneous X-ray and TeV $\\gamma$-ray data will allow us to constrain the jet parameters when modelling the emission processes.  Analysis of the Whipple data is ongoing, but attempts to reconstruct the source spectrum have been hampered by the effects of the decreased telescope efficiency during the observation period. The HEGRA collaboration measure a rather steep spectrum (spectral index $\\alpha=3.2\\pm0.3$) for observations prior to 2002, while the spectrum during the flaring period exhibits pronounced curvature and deviates significantly from the spectrum seen during the quiescent state \\citep{Horns02}. The majority of models for the EBL lead to predictions of a cut-off in the $\\gamma$-ray region beginning below $\\sim10\\U{TeV}$ for a source at $z=0.048$ (\\citet{Primack01}; but see \\citet{Vassiliev00} for a discussion of an EBL model which does not produce a distinct feature in the observed spectrum). Deviations from a pure power law have now been resolved in the spectra of both Markarian~421 \\citep{Krennrich01, Piron01, Aharonian02a} and Markarian~501 \\citep{Samuelson98, Aharonian99, Djannati99} and the spectrum of the most distant TeV blazar, H1426+428, is measured to be very steep ($\\alpha=3.50\\pm0.35\\mathrm{(stat)}\\pm0.05\\mathrm{(syst)}$) \\citep{Petry02}, but it is not yet clear whether these features are due to absorption on the EBL or are intrinsic to the sources. Clearly, further observations and spectral analysis of 1ES1959+650 may help to resolve this question."
    },
    "0212/astro-ph0212200_arXiv.txt": {
        "abstract": "{ While most chemical reactions in the interstellar medium take place in the gas phase, those occurring on the surfaces of dust grains play an essential role. Such surface reactions include the catalytic production of molecular hydrogen as well as more complex reaction networks producing ice mantles and various organic molecules. Chemical models based on rate equations including both gas phase and grain surface reactions have been used in order to simulate the formation of chemical complexity in interstellar clouds. For reactions in the gas phase and on large grains, rate equations, which are highly efficient to simulate, are an ideal tool. However, for small grains under low flux, the typical number of atoms or molecules of certain reactive species on a grain may go down to order one or less. In this case the discrete nature of the populations of reactive species as well as the fluctuations become dominant, thus the mean-field approximation on which the rate equations are based does not apply. Recently, a master equation approach, that provides a good description of chemical reactions on interstellar dust grains, was proposed. Here we present a related approach based on {\\it moment equations} that can be obtained from the master equation. These equations describe the time evolution of the moments of the distribution of the population of the various chemical species on the grain. An advantage of this approach is the fact that the production rates of molecular species are expressed directly in terms of these moments. Here we use the moment equations to calculate the rate of molecular hydrogen formation on small grains. It is shown that the moment equation approach is efficient in this case in which only a single reactive specie is involved. The set of equations for the case of two species is presented and the difficulties in implementing this approach for complex reaction networks involving multiple species are discussed. ",
        "introduction": "\\label{intro.}  The chemistry of interstellar clouds consists of reactions taking place in the gas phase as well as on the surfaces of dust grains \\citep{Hartquist1995}. Reactions that take place on dust grain surfaces include the formation of molecular hydrogen as well as reaction networks producing ice mantles and various organic molecules. The formation of H$_2$ on dust grain surfaces is a process of fundamental importance due to the fact that H$_2$ cannot form in the gas phase efficiently enough to account for its observed abundance \\citep{Gould1963,Hollenbach1970,Hollenbach1971a,Hollenbach1971b}. Molecular hydrogen plays a crucial role in gravitational collapse and star formation by absorbing energetic photons and radiating the energy in the infrared, to which the cloud is more transparent, thus cooling it and enabling further collapse. Furthermore, H$_2$ molecules are a necessary component for the initiation of chemical reaction networks that give rise to the chemical complexity observed in interstellar clouds. Therefore, the process of H$_2$ formation has attracted the attention of both theorists \\citep{Williams1968,Smoluchowski1981,Smoluchowski1983,Aronowitz1985,Duley1986,Pirronello1988,Sandford1993,Takahashi1999} and experimentalists \\citep{Brackmann1961,Schutte1976,Pirronello1997a,Pirronello1997b,Pirronello1999,Manico2001}.  The computational modeling of chemical reaction networks on dust grains in the interstellar medium is typically done using rate equation models \\citep{Pickles1977,d'Hendecourt1985,Brown1990,Brown1990b,Hasegawa1992,Hasegawa1993a,Hasegawa1993b,Caselli1993,Willacy1993,Shalabiea1994}. These models consist of coupled ordinary differential equations that provide the time derivatives of the densities of the species involved. Integration of these equations provides the time evolution of the densities. As long as the grains are not too small and the flux is not too low, the mean-field approximation applies. In this case, rate equations are an ideal tool for the simulation of surface reactions, due to their high computational efficiency. However, in the limit of very small grains under very low flux, rate equations are not always valid. This is because they take into account only average densities and ignore the fluctuations as well as the discrete nature of the populations of the atomic and molecular species \\citep{Tielens1995,Charnley1997,Caselli1998,Shalabiea1998,Stantcheva2001}. These features become significant in the limit of small grains and low incoming flux, typically encountered in diffuse interstellar clouds where hydrogen recombination as well as various other grain surface reactions take place. For example, as the number of H atoms on a grain fluctuates in the range of 0, 1 or 2, the H$_2$ formation rate cannot be obtained from the average number alone. This can be easily understood, since the recombination process requires at least two H atoms simultaneously on the surface.  Recently, a master equation approach was proposed, that is suitable for the simulation of chemical reactions on microscopic grains \\citep{Biham2001,Green2001,Biham2002}. It takes into account both the discrete nature of the H atoms as well as the fluctuations. In the case of hydrogen recombination, its dynamical variables are the probabilities $P_{\\rm H}(N_{\\rm H})$ that there are $N_{\\rm H}$ atoms on the grain at time $t$. The time derivatives $\\dot{P}_{\\rm H}(N_{\\rm H})$, $N_{\\rm H}=0, 1, 2, \\dots$ are expressed in terms of the adsorption, reaction and desorption terms. The master equation provides the time evolution of $P_{\\rm H}(N_{\\rm H})$, $N_{\\rm H}=0, 1, 2, \\dots$, from which the recombination rate can be calculated. The master equation approach has been applied to the study of reaction networks involving multiple species \\citep{Stantcheva2002,Stantcheva2002b}. Since the state of the system is given by the set of number densities of all species, the number of equations quickly increases with the number of reactive species. Therefore, suitable cutoffs should be imposed in order to keep the number of equations at a tractable level. A Monte Carlo approach based on the master equation was also proposed \\cite{Charnley2001}, and applied for reaction networks involving multiple species \\citep{Stantcheva2002b}.  In this paper we present a set of moment equations that can be derived from the master equation, and use it to calculate of the rate of molecular hydrogen formation on small grains. In the moment equations, the time derivatives of the moments $\\langle N_{\\rm H}^k \\rangle$, $k=1,2,\\dots$, of the distribution $P_{\\rm H}(N_{\\rm H})$ are expressed in terms of $\\langle N_{\\rm H} \\rangle$, $\\langle N_{\\rm H}^2 \\rangle$, $\\dots$, $\\langle N_{\\rm H}^{k+1} \\rangle$. To simulate the moment equations we impose a cutoff at a suitable value of $k$, giving rise to $k$ coupled linear differential equations. With such cutoffs, the moment equations can be used efficiently to calculate the formation rate of molecular hydrogen for both steady state and time dependent conditions. Furthermore, they provide a useful asymptotic expression for the rate of molecular hydrogen formation in the limit of very small grains. This expression can be used in order to integrate the H$_2$ production rate over the grain-size distribution in an interstellar cloud \\cite{Mathis1977,Draine1985,Draine1985,Mathis1990,Mathis1996,Weingartner2001,Cox2002}, and obtain the total production rate per unit volume of the cloud.  In the case of complex reaction networks involving multiple species, the number of moment equations quickly increases as the number of reactive species grows. Furthermore, the higher moments of the distribution are large and cannot be neglected. Thus, setting up the cutoffs in order to limit the number of equations is complicated and tends to introduce significant errors. As a result, for the case of multiple species we have not been able to develop an effective computational framework based on the moment equations. Here we present the set of moment equations for two species, such as hydrogen and oxygen, where the moments take the form $\\langle N_{\\rm H}^k N_{\\rm O}^{\\ell} \\rangle$, $k,l=0,1,2,\\dots$, and discuss the approximations that may be used in order to set suitable cutoffs in these equations. In spite of these difficulties, the moment equation approach has some advantages that justify further attempts to develop it for the case of multiple species. One advantage is that the moment equations are a direct generalization of the rate equations and resemble their structure. One may thus expect that they will be suitable for incorporation into models of interstellar chemistry based on rate equations. Another advantage is that in this approach, the moments, which quantify the production rate of molecules, are directly calculated unlike the case of the master equation where the calculation of moments requires further processing.   The paper is organized as follows. In Sect. 2 we consider reactions involving a single atomic specie, focusing on the case of hydrogen recombination. In Sect. 3 we consider reactions involving two atomic species, such as hydrogen and oxygen. In each of these Sections we briefly describe the rate equations and master equation for the system under study and then introduce and analyze the moment equations. The results are discussed and summarized in Sect. 4. ",
        "conclusions": "\\label{sec:Summary}  We have introduced a set of moment equations for the analysis of the formation of molecular hydrogen and other chemical reactions on dust grain surfaces in the interstellar medium. These equations are derived from the master equation that describes these processes. Like the master equation, the moment equations are exact even in the limit of small grains and low flux where the rate equations fail. They take into account the discrete nature of the populations of atomic and molecular species that participate in the reactions. While the master equation is expressed in terms of distributions such as $P_{\\rm H}(N_{\\rm H})$, $N_{\\rm H}=0,1,2 \\dots$, the moment equations are expressed in terms of the moments of these distributions.  We have shown that the moment equation approach is useful for simulations of molecular hydrogen formation. We expect that this approach will make it possible to efficiently calculate the total production rate of molecular hydrogen in diffuse clouds. The results will then be used in order to determine whether the molecular hydrogen formation process studied here is efficient enough to account for the observed abundance of H$_2$. The input required for such calculations includes the surface parameters (obtained from laboratory experiments), the grain size distribution, the density and temperature of the H gas, the grain temperature as well as the dissociation rate of H$_2$ molecules in the gas phase. Since the formation rate of molecular hydrogen is highly sensitive to the grain temperature and size distribution, the accuracy of the results may be limited by the quality of the observational data.  For more complex reaction networks involving multiple species, setting the cutoffs in the moment equations requires certain approximations that introduce significant errors. Therefore, for system that consist of more than a single reactive specie, we currently do not have a useful computational framework based on the moment equations. Potential advantages of the moment equations, justifying further attempts to further develop this approach, include the fact that they are similar in structure to the rate equations, possibly making it easier to incorporate them into models of interstellar chemistry. Moreover, they directly provide the first and second moments, which quantify the production rates of the molecular species involved. This is unlike the master equation in which the moments are calculated from the distribution in a post-processing stage."
    },
    "0212/astro-ph0212036_arXiv.txt": {
        "abstract": "{ Our VLA observations of the XMM-Newton/Chandra 13hr deep survey field (see Page et al., this proceedings) result in one of the two deepest 1.4GHz radio maps ever made.  Within the $15'$ radius field covered by the deep X-ray data (0.19 deg$^2$), a total of 556 radio sources are detected, down to a $4\\sigma$ flux limit of $28\\mu$Jy.  Of the 214 Chandra sources, 55 have radio counterparts.  The sub-arcsecond accuracy of the VLA and Chandra positions enable us to determine with high confidence the sources common to both surveys. Here we present the relationship between the X-ray and radio source populations at the faintest radio flux limits yet probed by such a study.  We discuss how the X-ray/radio relationship differs as a function of optical morphology, ie between unresolved `stellar' objects and well resolved galaxies. We then discuss the origin of the X-ray and radio emission, ie AGN, starburst or a mixture of both, in these two classes of object. ",
        "introduction": "Our Chandra/XMM-Newton X-ray survey of the 13hr field ($13^{\\rm h} 34^{\\rm m} 37^{\\rm s} +37^\\circ 54' 33''$, J2000) has been undertaken with the aim of understanding the physics of the X-ray sources around the knee of the source counts, ie, those sources which produce the bulk of the X-ray background radiation.  To complement our initial X-ray survey, we have made a deep VLA 20cm radio observation of the field, both as a survey in its own right, but also to investigate the overlap between the faint X-ray and faint radio source populations. Are they entirely different populations or different manifestations of the same object?  If the populations overlap, what are the underlying physics and emission mechanisms in these sources?  How does the emission mechanism depend upon host galaxy morphology?  (see also, eg, Bauer et al.~2002).  \\subsection{The Data}  The 13hr field was originally a deep ROSAT survey field (100ks with the PSPC, M$^{\\rm c}$Hardy et al.~1998), and was therefore chosen to be at high galactic latitude in a region of extremely low extinction, $N_H = 6.5 \\times 10^{19} {\\rm cm}^{-2}$.  The 13hr field has now been observed for 200ks with XMM-Newton (Mason et al., in preparation) as part of the Guaranteed Time programme, and with a mosaic of four 30ks Chandra ACIS-I pointings (M$^{\\rm c}$Hardy et al.~2002) covering 0.19 deg$^2$, reaching $1.3\\times 10^{-15}\\,{\\rm erg\\,cm^{-2}\\,s}^{-1}$.  A wealth of supporting data is available for this field, including a deep Subaru SuprimeCam optical image to $R\\sim 27$ and UKIRT/UFTI near infrared images of the bright X-ray galaxies.  This field will also be observed at mid to far infrared wavelengths as part of the SIRTF Guaranteed Time programme, and at 1mm with BOLOCAM.  Since the 13hr field is such a well studied field, it was also chosen for an extremely deep VLA survey (over 40hrs at 20cm, mostly with the A-array, plus a further 10hrs at 6cm) reaching $7\\mu$Jy rms, making it one of the three deepest radio surveys ever made (see also Richards et al.~1998; Ivison et al.~2002).  To study the brightest sources at higher resolution, 18 days of MERLIN 20cm observations were also performed.   \\subsection{Luminosities}  \\begin{figure} \\centerline{ \\begin{minipage}{68mm} \\resizebox{\\hsize}{!} {\\includegraphics[angle=270]{subMvlogLr_kfg.ps}} \\end{minipage}} \\caption{Log radio luminosity {\\it vs} $M_R$ for our sub-sample of radio sources with spectroscopic redshifts ($R<20.5$, 77\\% complete), showing broad line sources, ie, QSOs and BLRGs (filled stars), narrow line galaxies (open circles), and absorption line galaxies (filled squares).  The horizontal line at $L_R = 10^{25.6}$ denotes the conventional distinction between FR II and FR I sources.} \\label{lum} \\end{figure}  Since the VLA survey reaches to such faint fluxes, it is interesting to look at the typical luminosities for objects in our sample. Fig.~\\ref{lum} shows the radio luminosity {\\it vs} optical absolute magnitude for the sub-sample of sources with $R<20.5$ for which redshifts are available.  Of these, the majority of radio sources have $z<1$, and are of relatively low luminosity.  For the fainter sources, if we assume $z\\sim 1.5$, then a $50\\mu$Jy source has a luminosity of $3\\times 10^{23}\\,{\\rm WHz}^{-1}$, which is at the upper limit for `normal' galaxies; an X-ray source of $10^{-15}\\,{\\rm erg\\,cm^{-2}\\,s}^{-1}$ has $L_X = 6 \\times 10^{42}\\,{\\rm erg\\,s}^{-1}$, at the upper limit for starburst galaxies.  We note also that almost all the radio sources in our sample are smaller than 50kpc in size, ie, the size of a galaxy or smaller, consistent with a starburst origin for much of the emission, therefore we are not detecting the equivalent of classical double-lobed radio galaxies.   \\subsection{Stellarity}  The optical counterparts to the X-ray and radio sources in our samples are found from the $R\\sim27$ Subaru image of the field.  The {\\sc SExtractor} source extraction package (Bertin \\& Arnouts 1996) is used to create the optical catalogue, containing photometric and morphological information.  The `stellarity' parameter denotes how resolved an optical source is: broadly speaking, stars have stellarity $S = 1$ and galaxies $S = 0$, but a crude separation can be made at $S = 0.5$.  {\\sc SExtractor} has difficulty in classifying the morphology for objects fainter than $R\\sim 25$ in our data.  In Fig.~\\ref{stellarity}, we plot the stellarity of the optical counterparts against $R$--band magnitude, for both the X-ray and radio samples, which clearly shows that radio counterparts are more galaxy-like than the X-ray counterparts, as would perhaps be expected due to orientation effects in the context of the Unified Model of AGN.  \\begin{figure} \\centerline{ \\begin{minipage}{60mm} \\resizebox{\\hsize}{!} {\\includegraphics[angle=270]{stellar_chids.ps}} \\resizebox{\\hsize}{!} {\\includegraphics[angle=270]{radio_stellarity_rmag.ps}} \\end{minipage}} \\caption{The stellarity of the optical counterparts against their $R$--band magnitude for both our X-ray (upper) and radio (lower) samples.  } \\label{stellarity} \\end{figure}   \\subsection{Optical magnitude distributions}  \\begin{figure} \\centerline{ \\begin{minipage}{70mm} \\resizebox{\\hsize}{!} {\\includegraphics[angle=270]{kfg_rmag.eps}} \\resizebox{\\hsize}{!} {\\includegraphics[angle=270]{radio_rmag.eps}} \\end{minipage}} \\caption{The optical magnitude distribution for the X-ray (upper) and radio (lower) samples.  The shaded areas indicate those sources with spectroscopic identifications and redshifts.  The drop-off in sources fainter than $R>25$ is due to catalogue incompleteness and is not a real effect.  The bin at $R=27$ in each case reflects the sources for which an optical counterpart is not visible on the Subaru image.} \\label{optmag} \\end{figure}  In Fig.~\\ref{optmag}, the optical magnitude distributions of the X-ray and radio samples are plotted.  The Chandra sample (upper panel) has a strong peak at $23<R<24$, with very few counterparts with $R>25$.  The radio sample (lower panel) has a fairly flat magnitude distribution to $R\\sim 25$, and only falls off at $R>25$ due to optical catalogue incompleteness.  The higher fraction of optically faint sources means that more remain unidentified at present than in the X-ray sample. The broad range of optical magnitudes implies no correlation between radio or X-ray flux and host galaxy brightness. ",
        "conclusions": "From our VLA data, we obtain radio spectral and morphological information which, in combination with Chandra and XMM X-ray data, enable us to show the following results:  \\begin{itemize}  \\item 55/214 Chandra sources are faint ($>$28$\\mu$Jy peak flux) radio sources.  \\item Faint X-ray and faint radio emission is detected from a variety of sources, including relatively `normal' galaxies and QSOs.  \\item Wide ranges of $L_X/L_{\\rm opt}$ and $L_{\\rm rad}/L_{\\rm opt}$ imply a wide variety of emission mechanisms.  \\item At least one third of X-ray/radio matches are normal broad line AGN.  \\item The two optically brightest absorption line galaxies are almost certainly BL Lacs.  \\item The three optically brightest emission line galaxies are starbursts.  \\end{itemize}"
    },
    "0212/astro-ph0212085_arXiv.txt": {
        "abstract": "We study the mean environments of galaxies in the Sloan Digital Sky Survey as a function of rest-frame luminosity and color. Overdensities in galaxy number are estimated in $8\\,h^{-1}~\\mathrm{Mpc}$ and $1\\,h^{-1}~\\mathrm{Mpc}$ spheres centered on $125,000$ galaxies taken from the SDSS spectroscopic sample.  We find that, at constant color, overdensity is independent of luminosity for galaxies with the blue colors of spirals.  This suggests that, at fixed star-formation history, spiral-galaxy mass is a very weak function of environment.  Overdensity does depend on luminosity for galaxies with the red colors of early types; both low-luminosity and high-luminosity red galaxies are found to be in highly overdense regions. ",
        "introduction": "Elliptical and lenticular galaxies are over-represented in massive nearby galaxy clusters relative to the field \\citep{dressler80a, postman84a}.  Elliptical galaxies also tend to be redder, more luminous, more metal-rich, more gas-poor, and older in stellar population than spirals and irregulars \\citep[\\eg,][]{tammann79a, kormendy89a, roberts94a}.  Indeed, it has also been found that color, luminosity, surface-brightness, gas content, stellar population age, and star-formation rate are all correlated with the overdensity of the galaxy environment \\citep[\\eg,][]{kennicutt83a, balogh01a, martinez02a, lewis02a, blanton02d, gomez02q}.  Along the same lines, studies of the clustering of galaxies have found different clustering amplitudes for galaxies of different types, colors, and luminosities \\citep[\\eg,][]{davis76a, mo94a, park94a, norberg02a, zehavi02a}.  What is not understood is which of the relationships with environment are causal and which are just a by-product of other, more fundamental correlations.  Conventional cosmological theories posit that galaxies reside inside dark matter concentrations or ``halos'' that grow from small random fluctuations at early times.  The most overdense fluctuations will collapse first; in a gaussian random field, these preferentially reside within overdensities on larger scales, implying a correlation of halo mass with environment is expected \\citep[\\eg,][]{mo96a, lemson99a}.  The relationships between the properties of a halo and the properties of the galaxy or galaxies it contains are not fully understood, so the conventional theories do not currently make strong predictions for the dependence of galaxy number overdensity on observable galaxy properties, although most studies suggest that luminosity and color will be related to overdensity \\citep[eg,][]{kauffmann97a, kauffmann99a, benson00a}.  Whether or not this dark-matter halo picture ends up being useful or correct, and whatever turn out to be the important physical processes for making galaxies, the investigations started in this \\textsl{Letter} will place important constraints on galaxy formation and evolution.  The Sloan Digital Sky Survey (SDSS) is the best available data set for investigation of these relationships because of its sample size, high signal-to-noise imaging, sky coverage, and complete spectroscopy \\citep[\\eg,][]{york00a}.  Indeed, the SDSS has already made some relevant measurements, including the dependence of clustering on luminosity and color \\citep{zehavi02a}, the star-formation as a function of environment \\citep{gomez02q}, and the mean red-galaxy spectrum as a function of environment \\citep{eisenstein02a}.  In this \\textsl{Letter,} we investigate the mean galaxy number overdensities around galaxies of different colors and luminosities.  In what follows, a cosmological world model with $(\\Omega_\\mathrm{M},\\Omega_\\mathrm{\\Lambda})=(0.3,0.7)$ is adopted, and the Hubble constant is parameterized $H_0=100\\,h~\\mathrm{km\\,s^{-1}\\,Mpc^{-1}}$, for the purposes of calculating distances and volumes \\citep[\\eg,][]{hogg99cosm}. ",
        "conclusions": "We have shown that blue galaxies, \\ie, galaxies bluer than the red sequence of old stellar populations, exhibit no correlation between their luminosities and the overdensity of their environments, at fixed color.  The class of blue galaxies includes most spirals and irregulars.  On Fig~\\ref{fig:2d} a line of (roughly) constant stellar mass is shown, derived by converting relations derived in the $B$ and $R$ bands \\citep{bell01b}; the conversion is straightforward because $B$ and $R$ are very similar to $^{0.1}g$ and $^{0.1}i$.  The line of constant mass is close to a line of constant luminosity.  Similarly, a line of constant star-formation history is close to a line of constant color, with a small slope arising from the luminosity--metallicity relation \\citep[\\eg,][]{vila-costas92a, zaritsky94a, ryder95a}. Fig~\\ref{fig:2d} shows that, for spiral galaxies, mean overdensity is monotonically related to star-formation history.  That there is a strong relationship between star-formation history and environment has been suggested before, both theoretically \\citep[eg,][]{kauffmann99a, benson00a} and observationally \\citep{balogh01a, gomez02q}.  What is more remarkable is that, for galaxies with the colors of spirals, at fixed star-formation history, overdensity does not appear to depend on stellar mass at all.  The reddest galaxies show strong trends in overdensity at both the luminous and faint extremes.  The former trend indicates that the very most luminous galaxies lie in the very largest overdensities; probably this is related to the fact that giant galaxies lie near the centers of clusters.  The latter trend suggests that though high-density regions are expected to be rich in dwarf galaxies, those dwarfs will, by and large, be much redder than typical galaxies of their luminosity (\\citealt{norberg02a}; Zehavi et al, in preparation).  That red galaxies have a minimum in their mean overdensity at luminosities near $L^\\ast$ could be partly due to the specific mixture of galaxy types at those luminosities; there is a large overlap with the distribution of spirals, and the red population contains a large number of Sa galaxies at those luminosities \\citep[\\eg,][]{strateva01q}.  At one point it was thought that dwarf galaxies might ``fill the voids,'' making them an underdense population in the mean.  Our results rule this out, at least for dwarf galaxies selected by color and luminosity in the ranges considered here."
    },
    "0212/astro-ph0212566_arXiv.txt": {
        "abstract": "New observations of Cosmic Microwave Background Anisotropies, Supernovae luminosity distances and Galaxy Clustering are sharpening our knowledge about dark energy. Here we present the latest constraints. ",
        "introduction": "The discovery of a possible accelerated expansion of our present universe from type Ia supernovae (\\cite{super1}), is perhaps the most remarkable cosmological finding of recent years. Furthermore, the flatness of the universe ($\\Omega_{tot}=1$) determined by Cosmic Microwave Background observations (\\cite{flat}), togheter with the low matter density ($\\Omega_{matter}<0.4$) inferred from Large Scale Structure (\\cite{lowmat}) are suggesting the presence of a cosmological constant with high statistical significance. However, the CMB+LSS result relies on the assumption of a particular class of models, based on adiabatic primordial fluctuations, cold dark matter and a cosmological constant as dark energy component . In the following we will refer to this class of model as $\\Lambda$-Cold Dark Matter ($\\Lambda$-CDM).  This weak point, shared by most of the current studies, should not be overlooked: it might be possible that a different solution to the dark energy scenario than a cosmological constant can affect the CMB+LSS determination. It is therefore timely to investigate if the actual CMB data is in complete agreement with the $\\Lambda$-CDM scenario or if we are losing relevant scientific informations by restricting the current analysis to a subset of models.  Here first we check to what extent modifications to the standard $\\Lambda$-CDM scenario are {\\it needed} by current CMB observations with a model-independent analysis obtained fitting the actual data with a phenomenological function and characterizing the observed multiple peaks through a Monte Carlo Markov Chain (MCMC) algorithm, which allows us to investigate a large number of parameter simultaneously ($15$ in our case). We found a very good agreement between the position, relative amplitude and width of the peaks obtained through the model independent approach with the same features expected in a $4$-parameters model template of $\\Lambda$-CDM spectra. Second, since the $\\Lambda$-CDM is a good fit to the CMB data, we then move to other possible candidates for the dark energy component to see what kind of constraint we can obtain. The common characteristic of alternative scenarios to a cosmological constant, like quintessence or domain walls, is that their equations of state, $w_{dark}=p/\\rho$, might differ from the value for a cosmological constant, $w_{\\Lambda}=-1$. Observationally finding $w_{dark}$ different from $-1$ will therefore be a success for the alternative scenarios. We found that the present CMB and LSS data can put strong constraints on the values of $w_{dark}$ assumed as a constant. ",
        "conclusions": ""
    },
    "0212/astro-ph0212420_arXiv.txt": {
        "abstract": "We report on the properties of radio-selected galaxies within 30 very-rich Abell clusters with $z \\la\\,0.25$.  The radio, optical, and x-ray data for these clusters were presented in Paper I \\citep{mor00}. These radio data sample the ultra-faint ($L_{1.4} \\ge\\,2\\times 10^{22}$ W Hz$^{-1}$) radio galaxy population with $M_{\\mathrm{R}} \\le$\\,$-21$ using the well-known FIR/radio correlation to link the radio with ongoing star formation within individual cluster galaxies. Spectroscopic redshifts exist for $\\sim$\\,96$\\%$ of the optical identifications. These radio-selected galaxies reveal the `active' galaxy population (starburst and active galactic nuclei) within these rich cluster environments that can be identified regardless of their level of dust obscuration. These new radio data provide the largest sample to date of low-luminosity radio galaxies within rich cluster environments allowing an unbiased search for dusty starbursting galaxies. For all clusters in our sample, we are sensitive to star formation rates (M $\\ge$ 5M$_{\\sun}$) $\\ga$\\,5\\,M$_{\\sun}$yr$^{-1}$. We have found that the excess number of low-luminosity `starburst' radio-selected galaxies (SBRG) found by \\cite{owe99} in Abell 2125 is {\\it not} indicative of other rich clusters in our sample. The average fraction of SBRG is $\\langle\\,f_{\\mathrm{SBRG}}\\,\\rangle = 0.022\\pm0.003$. The A2125 fraction is $f_{\\mathrm{SBRG}} = 0.09\\pm0.03$ which is significantly different from the sample average at a $>$\\,99.99$\\%$ confidence level. Both A1278 and A1689 are slightly different from the rest of the sample at $\\sim$\\,90$\\%$ confidence level. The bimodal structure of both the x-ray brightness distribution and optical adaptively smoothed images of A1278 and A2125 suggests that ongoing cluster-cluster mergers may be enhancing this SBRG population. The A1689 excess low-luminosity (and high-luminosity) radio galaxy population may be due to interaction with the ICM. The mid-infrared ISOCAM results for A1689's radio galaxy population suggests that the radio emission for both low- and high-luminosity radio galaxies is AGN in origin except for one radio galaxy.  There is a significant spatial distribution difference between the low and high-luminosity (HLRG) radio-selected populations. The SBRG have a core radius of $0.40\\pm0.08$~Mpc which is $>$\\,$3\\times$ larger than the HLRG core radius. In addition, 48$\\%$ of the SBRGs have colors that are bluer than a typical Sab galaxy compared to 4$\\%$ for the HLRGs. The average absolute magnitude for the SBRG's is $\\langle\\,M_{\\mathrm{R}}\\,\\rangle =-21.93\\pm0.05$, while for the HLRG's it is $\\langle\\,M_{\\mathrm{R}}\\,\\rangle =-22.33\\pm0.07$, indicating that the SBRG are less optically luminous than their HLRG counterparts. The HLRGs seem to be a subclass of the cluster's massive red elliptical population, while the SBRGs have a projected radial distribution more like the blue spiral population. Our results indicate that most of the SBRGs are probably gas-rich disk galaxies undergoing $\\ga$\\,5\\,M$_{\\sun}$\\,yr$^{-1}$ of star-formation. ",
        "introduction": "Hierarchical clustering models predict that clusters of galaxies are assembled by a continuous coalescence of subclusters \\citep[e.g.,][]{whi76,evr90b,lac93,lac94,bal98a}. The ongoing mass accretion yields artifacts in the form of cluster asymmetry and significant amounts of subclustering. The remnants of structure formation are visible in current epoch clusters as substructure and cluster-cluster merging events. Estimates of detectable substructure in rich nearby clusters range from 30$\\%$-40$\\%$ \\citep[e.g.,][]{gel82, wes90}. The timescale for the accretion events is $\\sim$\\,1\\,Gyr for group-cluster encounters and $\\sim$\\,4\\,Gyr for cluster-cluster mergers\\citep{evr90b}, compared to $\\lesssim$0.3~Gyr for a typical starburst (SB)\\citep[e.g.,][]{ken98}. Such SB episodes show up as a cluster-wide enhancement of the radio-selected galaxy population. Therefore, while the substructure caused by large-scale-structure  formation will be visible for a significant fraction of a Hubble time, the cluster-wide enhanced radio activity, as in the case of A2125 \\citep{owe99}, will be rather brief.  The Butcher-Oemler (BO) effect \\citep{but84b} has been suggested as a link between galaxy evolution and cluster dynamics. The hierarchical models predict that distant clusters assemble from smaller subclumps at a much higher rate than similar mass nearby clusters. These smaller groups have a higher fraction of gas-rich, late-type galaxies, thereby providing a distant cluster with a population that would support significant amounts of induced star formation (SF). In fact, these late-type galaxies may be transformed into early-type galaxies by environmental dynamics of the cluster \\citep[e.g.,][]{oem97}.  The \\cite{cal97} spectral study of early-type galaxies in rich clusters with significant substructure shows that $\\sim$\\,15$\\%$ of these galaxies have signatures of current or ongoing SF. This suggests that cluster-group merger activity signified by the substructure may alter the star formation rate (SFR) of a cluster galaxy population. Induced SF may result from shocks caused by the collision between the intracluster medium of the cluster and that of the group. Alternatively, the main driver of this activity may be galaxy-galaxy collisions in regions where the velocity dispersion is low enough and the density is high enough for such encounters to be efficient. The different regions that the starburst population inhabits within the rich cluster gives us clues as to the triggering mechanism for this activity. Radio observations allow us to identify this population without being hampered by dust extinction or the K-correction.  This paper is the second in a series studying the ultra-faint radio-selected galaxy populations associated with rich clusters of galaxies as a function of redshift.  We analyze both the high and low luminosity radio-selected galaxy population within 30 very rich Abell clusters (richness $\\ge$2) with $z \\la 0.25$. The data used in our analysis was presented in paper I \\citep{mor00}.  Initially, \\cite{dwa99} conducted a VLA 20 cm wide-field continuum study of two very rich Abell clusters, A2125 and A2645, looking at the radio galaxy population down to $\\sim$2$\\times$10$^{22}$ W Hz$^{-1}$. The reason for choosing this lower radio luminosity limit is that \\cite{con91} had determined that in the local universe this limit has a statistically higher fraction of starburst-powered radio-selected galaxies and few galaxies whose radio emission is powered by AGNs. Both clusters are at the same redshift (z=0.25) and richness (R=4), but A2125 has a higher blue galaxy population ($f_{\\mathrm{B}}$ = 0.19) than A2645 ($f_{\\mathrm{B}}$ = 0.03). The radio galaxy population is also substantially different: A2125 has 26 radio detected galaxies whereas A2645 has 4 \\citep{owe99}.  Most of the radio galaxy population of A2125 is not confined to the core region of the cluster ($r <$ 400 kpc) but is distributed out to a radius of 2.5 Mpc. The radio-selected galaxies in the cluster core are red in color, with none of the blue Butcher-Oemler galaxies detected down to the radio luminosity limit. The origin for the radio emission is both AGN and massive ongoing SF. The latter occurs more frequently in the outer cluster region of A2125.  This result raises several interesting questions. Are we just seeing the infall of the field population at $z = 0.25$, or are we witnessing an enhancement in the activity level of the population due to the cluster environment?  Do all very rich clusters show enhanced activity in the form of starburst and AGNs, or only the ones at higher redshift?  To understand if this is a richness or a redshift effect, we have constructed a sample of 34 very rich clusters from 0.02 $\\la z \\la$ 0.41 to study how the radio properties of these clusters evolve (if at all) over the last $\\sim$5 Gyr.  VLA radio observations at 20\\,cm offer a wide field-of-view ($\\sim$30 arcmin), allowing us to sample out to 2.5 Mpc radius from the cluster core. Nearby clusters ($z \\leq 0.06$) were analysed using the NRAO VLA Sky Survey (NVSS) which provided the required linear coverage (5\\,Mpc diameter search area) and sensitivity.  At this projected distance from the cluster core we should detect any infalling starbursting field galaxies. Our detection limit of 2$\\times$10$^{22}$ W Hz$^{-1}$ ($H_0$ = 75\\,km\\,s$^{-1}$\\,Mpc$^{-1}$; $q_0 = 0.1$) yields a SFR(M $\\ge$ 5M$_{\\sun}$) sensitivity limit of $\\ga 5$ M$_{\\sun}$ yr$^{-1}$ assuming the SFR relation from \\cite{con92}.  In this paper, we analyze the radio properties of 30 of the richest, nearest Abell clusters from 0.02 $\\la$ z $\\la$ 0.25. These clusters were chosen to be as rich or richer than the four $z = 0.4$ clusters. We investigate the characteristics of the different radio galaxy populations and their relationship to the cluster properties, such as richness, core richness, compactness, and x-ray luminosity. In addition, the Butcher-Oemler-defined blue color and absolute optical luminosity of the radio-selected objects are also explored. ",
        "conclusions": "In this paper we have studied the radio-selected galaxy population of a subsample of very rich ($R \\ge 2$) Abell galaxy clusters with $z \\le 0.25$.  The weak anti-correlation between $f_{\\mathrm{RG}}$~and $N_{2.0}$~at the 90$\\%$ confidence level suggests possibly that the detection rate for all radio galaxies does not scale with the number of galaxies surveyed.  However, overall, we found no significant ($\\geq 2\\sigma = 95\\%$ confidence) correlation or anti-correlation of $f_{\\mathrm{RG}}$, $f_{\\mathrm{SBRG}}$, or $f_{\\mathrm{SBRG}}$ with cluster richness ($N_{2.0}$ or $N_{\\mathrm{0.5}}$), compactness, blue fraction, or x-ray.  There are only a few clusters that have radio galaxy fractions (SBRG and HLRG) that are inconsistent with the rest of the sample.  We find the following average radio galaxy fractions: $\\langle\\,f_{\\mathrm{SBRG}}\\,\\rangle = 0.022\\pm0.003$ and $\\langle\\,f_{\\mathrm{HLRG}}\\,\\rangle = 0.024\\pm0.003$.  The cluster with an excess at the $>~99.99\\%$ confidence level of SBRGs is A2125, with a $f_{\\mathrm{SBRG}} = 0.09\\pm0.03$. This is {\\it not} a richness-induced effect as we normalized by the number of galaxies sampled.  Two clusters with weak $\\sim 90\\%$ confidence level deviations from $\\langle\\,f_{\\mathrm{SBRG}}\\,\\rangle$ are A1278 at $0.05\\pm0.03$ and A1689, at $0.05\\pm0.03$.  The bimodal structure in the optical and X-ray for A1278 and A2125 suggests that a possible cluster-cluster merger may be driving this excess in SBRGs. The mechanism for the increased fraction of SBRGs in A1689 may be the result of the ISM of the galaxies being compressed by their passage through the ICM. All the SBRGs in A1689 are within a projected distance of 0.8\\,Mpc from cluster core.   The large fraction of HLRGs in A1689 ($0.09\\pm0.04$) and A1940 ($0.06\\pm0.03$) may also be a result of the cluster environment. The HLRGs close projected distance from the cluster center ($< 0.8$Mpc for both clusters) suggests that this population may be enhanced due to pressure confinement by the ICM. The MIR ISOCAM results for A1689's radio galaxy population suggest most are AGN in nature lends support to this idea.  The different spatial distribution between the SBRGs and the HLRGs is one of the most significant results of this paper. The core radius values of the SBRGs is $r_{\\mathrm{c}} = 0.40\\pm0.08$~Mpc. These are a factor $>$\\,$3\\times$ larger than the average value of $r_{\\mathrm{c}}= 0.12\\pm0.02$~Mpc for HLRGs. The large difference in the fitted core radii values of the SBRGs and the HLRGs indicates a strong difference exists between their representative populations. The HLRGs are probably a subclass of the cluster's massive {\\it red} elliptical population, while the SBRGs have a distribution more like the {\\it blue} spiral population.  The difference in $\\langle\\,M_{\\mathrm{R}}\\,\\rangle$ values between the SBRGs and the HLRGs indicates an absolute magnitude segregation between the two populations, with the higher optical luminosity galaxies belonging to the HLRGs.  This is in agreement with the core radius results for the SBRGs and HLRGs. Also the colors of the three radio populations suggest that the SBRGs have colors that are much bluer than the HLRGs.  These results indicate that a large fraction of the SBRGs are probably gas-rich disk galaxies with SFR $\\ga$5\\,M$_{\\sun}$\\,yr$^{-1}$. It is unclear as to what triggered the SF in the SBRGs for most clusters. It must be noted that contamination due to AGNs in the SBRGs has not been completely removed. Paper IV will cover the spectroscopic classifications of the radio-selected galaxies.  A larger, more morphologically diverse sample is currently being studied that contains more irregular rich clusters similar to A1278 and A2125. This will allow us to decouple the effects that cluster dynamics have on the radio properties of cluster galaxies."
    },
    "0212/astro-ph0212434_arXiv.txt": {
        "abstract": "A hot phase of the interstellar medium has now been detected and studied in several objects through its X-ray emission.  A proper assessment of its characteristics  is relevant for our understanding of several aspects of galaxy properties, from the large scale distribution of matter to the stellar and galaxian evolution, to the dynamics of systems and to the feeding of a central black hole. I will briefly summarize our current understanding of some of the main issues related to the hot gaseous component in galaxies that are fast evolving given the ever more striking and interesting details provided by the X-ray satellites currently operating. I hope to convince you that the  X-ray characteristics of the hot gas are quite complex, both in morphology and spectra, in a wide range of objects,  which should promote greater efforts in understanding the role played by this component in all galaxies. ",
        "introduction": "Because of the implications and relations that a hot phase of the interstellar/intergalactic medium has with several other galaxy-related issues, understanding its properties has become more relevant and intriguing. As just a few examples, gas is closely related to:  \\smallskip\\noindent $\\bullet$ galaxy structure.  Gas can be traced at large galactic radii and has been used to study and measure the total gravitational mass in early type galaxies (e.g. Fabricant \\& Gorenstein 1983) \\\\ $\\bullet$ evolution of stars, of galaxies and of larger systems.  Metal enriched material that is formed in stars gives a measure of stellar evolution; metal gradients indicate how far (how?) metals have traveled from their formation sites; metal concentrations in specific locations could indicate local stripping phenomena \\\\ $\\bullet$ the dynamics of systems.  At high energies different phenomena such as stripping and shocks could be discovered and studied \\\\ $\\bullet$ active nuclei, who are nourished by gas accretion.  When hot, the gas could indeed starve the nucleus rather than feed it, and explain the lack of activity or periodic cycles \\\\  \\smallskip\\noindent The main question is then to find the best candidates to efficiently study the properties of this phase of the InterStellarMedium (ISM).  \\begin{figure} \\unitlength1.0cm \\begin{picture}(11,6.0) \\thicklines \\put(0,-0.1){ \\psfig{file=trinchieri_fig1.ps,width=12cm,height=6cm} } \\end{picture} \\caption[]{From \\citeauthor{stricketal} (\\citeyear{stricketal}). Top panels: X-ray images in the 0.3-2.0 keV energy band from Chandra data for 4 galaxies from normal to actively forming stars.  Lower panels:  The same galaxies in the continuum-subtracted H$\\alpha$+[N{\\sc ii}] emission. The arrow at the bottom indicates the increasing IRAS $f_{60}/f_{100}$ ratio, which is used as a measure of the star formation (SF) intensity. The increasingly strong soft X-ray emission associated with increasing SF activity is evident. } \\label{strick} \\end{figure} ",
        "conclusions": "It is probably too early to draw conclusions from the evidence accumulated so far.  However, it is becoming clearer that the small scale gas morphology indicates that the ``early- type'' galaxy population is far from  ``quiescent''. Same can be said for ``cD'' galaxies in small/sparse groups.   This would indicate that there is ground for more detailed studies of the high energy emission in galaxies; these might reveal unsuspected phenomena like shocks that are crucial for our understanding of the evolution of these systems.  The multiwavelength approach is the only one that can provide a complete picture of all the forces at play, and the study of the gaseous components, though complex, has a very high potential for easy discovery of extreme, but perhaps not uncommon, phenomena."
    },
    "0212/astro-ph0212328_arXiv.txt": {
        "abstract": "Monte-Carlo computations for highly relativistic parallel shock particle acceleration are presented for upstream flow gamma factors, $\\Gamma=(1-V_{1}^{2}/c^{2})^{-0.5}$ with values between 5 and $10^{3}$. The results show that the spectral shape at the shock depends on whether or not the particle scattering is small angle with $\\delta \\theta < \\Gamma^{-1}$ or large angle, which is possible if $\\lambda > 2r_{g} \\Gamma^{2}$ where $\\lambda$ is the scattering mean free path along the field line and $r_{g}$ the gyroradius, these quantities being measured in the plasma flow frame. The large angle scattering case exhibits distinctive structure superimposed on the basic power-law spectrum, largely absent in the pitch angle case. Also, both cases yield an acceleration rate faster than estimated by the conventional, non-relativistic expression, $t_{acc}=[c/(V_{1}-V_{2})] [\\lambda_{1}/V_{1}+\\lambda_{2}/V_{2}]$ where '1' and '2' refer to upstream and downstream and $\\lambda$ is the mean free path. A $\\Gamma^{2}$ energy enhancement factor in the first shock crossing cycle and a significant energy multiplication in the subsequent shock cycles are also observed. The results may be important for our understanding of the production of very high energy cosmic rays and the high energy neutrino and gamma-ray output from Gamma Ray Bursts (GRB) and Active Galactic Nuclei (AGN). ",
        "introduction": "There are three distinct astrophysical situations where the bulk plasma flow is extremely relativistic with Lorentz factors $\\Gamma=(1-V^{2}/c^{2})^{-1/2} \\geq 10$. These are some Active Galactic Nuclei (AGN) jet and hot spot sites, Gamma Ray Burst (GRB) fireballs and pulsar polar winds. Within these flows relativistic shocks appear and diffusive particle acceleration takes place. In each case, there is evidence for energetic particle acceleration to high $\\Gamma$ factors but the upper limit to the possible energy attained becomes less certain with increasing bulk flow velocity. For AGN jets, $\\Gamma \\sim 10$ plausibly results in the known $10^{19}$ eV cosmic rays (Quenby and Lieu, 1989) via diffusive shock acceleration at the termination shocks. However, the more direct synchrotron/self-Compton modeling only reveals 10 GeV gamma-rays (e.g. Maraschi et al., 1992). For GRB fireballs, $\\Gamma\\sim 100-10^{3}$ appears to eventually produce gamma-rays at least up to 100 MeV (Fenimore et al., 1993) but at such flow speeds, it is not clear whether repeated shock crossings are possible and the predictions of $10^{19}$ eV neutrinos (Vietri, 1998), which would be direct evidence of a diffusive like process have yet to be verified. Similarly, pulsar winds could result in a $\\Gamma \\sim 10^{4}$ flow encountering the nebula envelope (Hoshino et al., 1992), but there is no evidence for more than TeV acceleration anywhere in the system. However Lucek and Bell (1994) show that combined motion across a near perpendicular shock with a non-uniform field and in the electromagnetic field of the nebula may produce much higher energies.\\\\ Past work has suggested that if the shocks are close to parallel so that an ${\\bf E}=0$ frame can be defined, diffusive shock acceleration is in principle possible and provided large angle scattering occurs in the plasma frame, each shock crossing cycle yields a $\\sim \\Gamma^{2}$ energy increase (Quenby and Lieu, 1989). This result has been exploited by Vietri (1995,1998) modeling GRB fireball particle shock acceleration. It is the purpose of this paper to present simulation results on particle  acceleration with $\\Gamma$ values extended up to $\\sim$ $10^{3}$. Here we concentrate on the  parallel shock configurations ($\\psi=0^{o}$). Our results, measured in the shock frame, will be compared  with the standard spectral prediction of non-relativistic theory, where the momentum spectrum is $n(p)\\propto p^{-\\alpha}$ with $\\alpha=(r+2)/(r-1)$, for shock compression ratio $r=V_{1}/V_{2}$, independent of scattering details, and $\\alpha=2$ for strong unmodified shocks with $r=4$. Analytically it has been shown that over a limited momentum range, the non-relativistic time scale for acceleration is $T=[c/(V_{1}-V_{2})][\\lambda_{1}/V_{1}+\\lambda_{2}/V_{2}]$ where, '1' and '2' refer to upstream and downstream and $\\lambda$ is the scattering mean free path. The shock jump solutions of Ballard and Heavens (1991) and the review of Kirk and Duffy (1999) show that, a range of shock frame $r$ values largely limited between 3 and 4, typically arise for oblique, relativistic shocks, whereas $r$=4 is the extreme relativistic hydrodynamic limit. Hence for simplicity we adopt $r$=4 but, in the companion paper dealing with oblique shocks we demonstrate the insensitivity of our chief conclusions on the exact value. For our parameters, $T=(8/3)(nr_{g,1}/c)$  where, $\\lambda=nr_{g}$ and $r_{g}$ is Larmor radius.\\\\ The first relativistic correction noticed to the non-relativistic theory, was the spectral flattening seen in parallel shocks (Kirk and Schneider, 1987a,b), followed by the discovery of a speed-up in the acceleration rate (Quenby and Lieu, 1989; Ellison et al., 1990). These results have subsequently been extended to non-parallel and non-linear sub-luminal shocks (Lieu, et al., 1994 and Bednarz and Ostrowski, 1996 who use $T=r_{g}/c$ as the unit of acceleration time), where there is large angle scattering or pitch angle diffusion. Most recently, much interest has been focused on differences occurring between small angle and large angle scattering models for the test particle-turbulence interactions in the fluid frame. Gallant and Achterberg (1999) suggest that for a field model consisting either of randomly orientated uniform field cells or a uniform field, the scattering is limited to $\\delta \\theta < 1/\\Gamma$ in the upstream fluid frame. Generally, such a model yields a test particle differential spectral slope $\\sim -2.2$ (Baring, 1999). Furthermore, on this model, Gallant and Achterberg (1999) provided an analytical demonstration that in the high $\\Gamma$ limit, only the first crossing cycle exhibited a $\\Gamma^{2}$ energy increase, while in all other crossings the energy is only doubled and hence one may expect the relativistic speed-up to be limited to this crossing. The computations of Baring (1999) also showed a much reduced energy enhancement in subsequent cycles for the small angle scattering case. Achterberg et al. (2001) extend the conclusions of Gallant and Achterberg (1999) to a wider variety of upstream field models while, Kirk et al. (2000) employ an eigenfunction method to deal with the anisotropies in diffusion when the scattering depends on the angle with the shock normal. A recent general review of particle acceleration and relativistic shocks is given by Kirk and Duffy (1999). However, the detailed contrast between the effects of various particle scattering models and the behaviour of these models at $\\Gamma$ factors $\\sim 10^{3}$ remain to be elucidated, both for parallel and quasi-perpendicular shocks. Hence this paper starts a series of investigations in the ultra-relativistic limit with a consideration of parallel shocks. In section 2 we present the numerical method used, section 3 presents the results while in section 4 conclusions and consequences are given. A further paper seeks to extend this work to non-parallel shocks (Meli and Quenby, 2002b). ",
        "conclusions": "To what extent is first order Fermi acceleration universally applicable in shocked, turbulent, MHD flows ? Aspects of the problem still needing elucidation include the spectrum produced, the spectral index, the angular distribution of the particles and non-linear effects in relativistic flows.\\\\ We have presented a detailed Monte Carlo approach for the highly relativistic diffusive acceleration in parallel shocks under the test particle approximation, extending previous studies to more extreme values of the upstream flow $\\Gamma$ and carefully distinguishing the effects of different assumptions concerning the particle scattering. Comparison is made with results under non-relativistic conditions. It is found that the spectral shape in the very relativistic flow regime depends on the scatter model (large angle or pitch angle diffusion/small angle scattering), with the former model exhibiting a structure in the spectrum at the shock related to the cycle of acceleration, while the later one showing a smoother shape. One general indication which supports the validity of  our computational approach is that the spectral index becomes flatter as the flow becomes more relativistic, a result to be compared with the conclusions of Kirk and Schneider (1987a,b).\\\\ Extreme anisotropy in the shock frame for particles emerging downstream irrespective of scattering model is observed, due to the beaming of the transmitted flux. High anisotropy also occurs in the flux emerging upstream in the small angle scattering case, due to the lack of time to isotropize the particles but in the large angle scattering model there is a large enough probability of scatter  before shock re-entry for near isotropy to set in. This means that for any given acceleration cycle, the spectrum is much flatter in the large angle case (plateaus) because the gain per cycle can be greater. However, evidently because of an enhanced loss downstream of the more scattered particles of the large angle scattering case, the average spectral slope for the two mechanisms is similar. \\\\ Most important, a significant acceleration rate increase ('speed up' effect) has been found for both large angle and pitch angle diffusion. This must be seen in conjunction with a 'gamma-squared' energy boosting of the particle noted in the first shock cycle, which is in accordance with theoretical indications (e.g. Gallant  and Achterberg, 1999) and past simulations (e.g. Lieu and Quenby, 1989). There is a gain by a significant factor in subsequent crossings which is $\\sim \\Gamma$ or even more for several cycles. Regarding the acceleration time decrease, computations show that the comparison of the simulation (experimental) acceleration time with the analytic non-relativistic time constant, indicates a decrease of about 5 with increasing flow gammas (10-$10^{3}$). This speed-up is clearly directly related to the boosting of the energy gain.\\\\ In the theory of GRB production (Meszaros and  Rees, 1993)  % by cosmic fireballs, the $\\Gamma \\sim 10^{3}$  flows are expected to produce some shock acceleration of electrons and protons. Vietri (1995) suggests protons of $\\gamma \\leq 10^{11}$ produce photons of 300 GeV by proton synchrotron radiation and neutrinos of $10^{19}$ eV via photopion production and subsequent cascading. This requires only two shock acceleration cycles at flow $\\Gamma\\sim 10^{3}$, a suggestion well supported by the results obtained here for both scattering models. However, at these high $\\Gamma$ values, extreme 'billard ball' scattering is required over distances much less than $\\lambda$ in the large angle case. What finally limits the proton maximum energy is escape from the shock turbulence region, as the Larmor radius of the accelerated particles increases in size. Despite the more favourable radiative process, electrons are not able to produce gamma-ray synchrotron energies higher than those due to these protons because the competing synchrotron energy loss severely limits the diffusive shock gain (e.g. Gallant et al., 2000) so, the net result is a similar upper photon cutoff. Because the accelerated particle spectra at the shock as here are not necessarily smooth, it might be expected that a detailed experimental study of radiation from this region of a GRB may reveal a related structure. In turn, such a study may reveal new insight into the magnetic field regime of the GRB environment.\\\\ The suggestion by Vietri (1995) that the efficient production in GRB of $10^{20}$ eV protons constitutes a cosmic ray source, at least in previous epochs, is also supported. This argument depends on the similarity of the burst gamma ray energy output and very high energy cosmic ray energy flux (Vietri, 1995; Waxman, 1995).\\\\ Finally, we may conclude that first order Fermi acceleration in the sense of additive energy change over a number of cycles of magnetic discontinuity crossing, appears to be applicable to $\\Gamma\\sim 10^{3}$ flows where, there is no test particle reflection at the discontinuity. However, the large gradients in particle distribution  function appearing in the simulations argue against a spatial diffusion approximation description of the phenomenon.\\\\  {\\noindent \\bf Acknowledgments}\\\\  We wish to thank Prof. Drury and the unknown referee for their valuable comments and suggestions regarding this work."
    },
    "0212/astro-ph0212058_arXiv.txt": {
        "abstract": "We discuss new developments of interstellar chemistry, with particular emphasis on the carbon chemistry. We confirm that carbon chains and cycles are ubiquitous in the ISM and closely chemically related to each other, and to carbon. Investigation of the carbon budget in shielded and UV illuminated gas shows that the inventory of interstellar molecules is not complete and more complex molecules with 4 or more carbon atoms must be present. Finally we discuss the consequences for the evolution of clouds and conclude that the ubiquitous presence of carbon chains and cycles is not  a necessary consequence of a very young age for interstellar clouds. ",
        "introduction": "\\label{intro}  Nearly 120  molecules have now been  observed in various sources in the interstellar medium. Emission or absorption lines from these molecules carry important information on the physical conditions and the chemical properties of the various sources where they have been recognized. However, only a very limited number of molecules is used for that purpose, while much more information is available. The use of molecular lines detected with high resolution spectroscopy for studying the interstellar medium, and the star formation processes, are still largely unexplored. High resolution spectroscopy, either using heterodyne techniques in the radio domain, or with high dispersive power spectrometers, yields not only the line intensities but information on the gas dynamics through the analysis of line profiles. Many lines can be analysed together for studying  physical conditions. Different approaches can be used to solve the radiative transfer equation, from the simplest LTE model to sophisticated multidimensional models. Also, other line properties can be exploited to get information on the magnetic field, either directly from the Zeeman effect on paramagnetic molecules, or indirectly with other techniques (e.g.  \\cite{Houde:03}). And finally, the molecular abundances are the basic information for the chemistry. The question which arises then is how to select the best molecule and the best lines to get useful  information.  All interstellar molecules can not be detected in the various environments of the ISM. This evolution of the emerging spectrum, hence of the chemical properties, is tightly connected to the physical conditions and environment of the source. With the advent of more sensitive receivers and better mapping techniques, it is now possible to study extensively the chemical properties of specific sources. For example  \\cite*{Tur:00} have performed detailed observational studies and modeling of a large set of interstellar molecules in cloud cores. We present in this paper new studies  of interstellar chemistry devoted to carbon chemistry and the condensation processes on dust grains, which make use of molecular line observations. ",
        "conclusions": "In all interstellar environments, dark and diffuse clouds, and PDRs, carbon chains and cycles are closely related to each other, and to atomic carbon. They are present in large abundances in the moderately dense molecular gas but avoid the darkest and densest cores. This conclusion is valid for the pure carbon chains and cycles. Other carbon species, like cyanopolyynes are found in high density regions. Though their formation mechanism also involves carbon, the difference between nitrogen and carbon chemistry may contribute to the different spatial distribution. The difference of incident radiation between diffuse or UV illuminated gas on one side, and dark clouds on the other side, affects somewhat the relative abundance of carbon molecules in permitting the growth of larger chains (compare the abundance of C$_4$H in TMC-1 and in the horsehead nebula) but the chemistry looks nevertheless very similar. Another major difference between diffuse or UV illuminated gas on one side, and dark clouds on the other side, is the condensation of abundant gas phase molecules on dust grains. The main impact of condensation on the chemistry is the enhanced abundance of molecular ions in depleted region. The reason for this is that the destruction channels for molecular ions are not only dissociative recombination reactions with electrons, but also reactions with abundant gas phase molecules. The depletion of CO on dust grains therefore leads to lower destruction rates for molecular ions and larger abundances. While the composition of molecular ice depends mainly on the total extinction, marked chemical differences are found in lines of sight with similar column densities. Also, measured abundances of carbon chains and cycles are often larger than steady state predictions. The usual explanation for understanding the large abundances, is to invoke the so-called \"early time chemistry\", in which clouds are seen in an early evolutionary stage since their formation, in which the available carbon, initially ionized,  has not be completely transformed into CO. However, several arguments show  that  \"early time chemistry\" cannot be the sole explanation.  Early time chemistry is efficient because the formation of molecular hydrogen and the conversion of carbon to CO takes a very long time ($\\geq 10^6$ years) when starting from atomic gas. It is thus tempting to use molecular abundances in dark clouds as a chemical clock. However the same arguments leading to doubt on ``early time '' chemistry, are also against this ```chemical clock '' idea for dense cores.  i) Since the same molecules, with similar abundances, are seen in many different clouds, at different evolutionary stages, such as PDRs, pre-stellar clouds, clouds having already formed stars, how could we be viewing ``early time'' chemistry in all cases ? For example, clouds having already formed stars and clouds associated with PDRs cannot be presented as ``young'' since they have already had time to collapse and form stars, and for PDRs the newborn stars have had time to disrupt their parent clouds. Even if the evolution is not uniform across a cloud (chemical time scales depend on the density, so that more diffuse regions evolve at a slow pace compared to dense cores) the requirement of ``early time'', i.e. a cloud age lower than  10$^5$ years is very strong compared to the typical time scale for cloud formation and evolution towards forming stars ($\\geq 10^6$ years, \\cite{hartmann:01}).  ii) A second argument against ``early time'' chemistry is the heavy dependence of the predicted chemical abundances on the initial conditions \\cite*{Lee:96} while observations show well defined chemical properties.  iii) Finally, as stated above, molecular ices are commonly detected in dark clouds which indicate the gas has had time to evolve, and built gas phase and solid phase molecules (\\cite{Roberts:02}).  Other processes are known to strongly affect the chemistry, and may favor the presence of carbon and carbon chains in cloud envelopes. The propagation of UV radiation depends on the grain optical properties. \\cite*{Casu:01} have shown that grain growth in dark clouds affects the grain optical properties hence the propagation of UV radiations and the photodissociation rates in dense gas. This effect leads to higher abundances of neutral and ionized carbon compared to models using the grain properties valid for diffuse gas. It must be remembered that photons play a very important role in interstellar chemistry, even in dark clouds since the interaction of cosmic rays with hydrogen leads to the production of UV photons deep into the clouds. In addition, the chemical equilibrium for moderately dense gas  can be carbon rich, the HIP solution (\\cite{LB:95}) even without any incident UV photons. These last two explanations rely on microscopic processes and choice of elemental abundances. But the structure and dynamical state of interstellar clouds also affects the chemistry. Clumpy cloud models, such as those of \\cite*{spaans:97} or \\cite*{storzer:96} can explain the widespread distribution of carbon by the accumulation of many clump surfaces, rich in carbon, in telescope beams. Since UV radiation can penetrate deeper into clumpy clouds than in homogeneous structures, each clump surface behaves as a mini-PDR in this context. Another idea is to include the effect of turbulence in homogeneous clouds. The main effect of turbulence is to increase the  diffusion between the cloud surface and its interior. A small level of turbulent diffusion, compatible with the observations, can maintain higher abundances of atomic carbon and atomic hydrogen in cloud envelopes than in steady state models, in agreement with observations of atomic hydrogen in dark clouds (\\cite{Willacy:02}).  It is likely that all additional physical mechanisms described above, both microscopic (change of optical properties of dust grains, abundances) and macroscopic (turbulence, chemistry) and possibly others still to uncover, contribute significantly to the chemistry. There are  so many unknowns in the physical properties and chemistry of dark clouds, that deriving the age of clouds from their chemical abundances is not reliable. Clues from the age of embedded stars, if any, the geometry, the environment etc.  must be used to retrace the history of interstellar structures. To progress, the whole evolution of interstellar clouds, from  their formation in the warm neutral gas, to their concentration in molecular gas, and finally the birth of new stars, must be better understood than it is now."
    },
    "0212/astro-ph0212572_arXiv.txt": {
        "abstract": "Cold, dense gas clouds have been proposed as a major component of the Galactic dark matter; such clouds can be revealed by the gamma-ray emission which arises from cosmic-rays interacting with the gas. If this dark matter component is clustered then highly luminous GeV sources result, preferentially at low to mid Galactic latitudes where they lie within the cosmic-ray disk. The predicted emission for such clusters is steady, continuum emission with peak power emerging at several hundred~MeV.  These sources would not have obvious counterparts at other wavelengths and are thus of interest in connection with the UnIDentified (UID) EGRET sources. Here we present a Monte Carlo simulation of the gamma-ray source population due to cold gas clouds, assuming a Cold-Dark-Matter-like mass spectrum for the clustering. We find that $\\sim280$ EGRET sources are predicted by this model, with a median Galactic latitude of $12^\\circ$ for the population, and a median angular size of $2.2^\\circ$. The latitude and size distributions are consistent with the UID EGRET source data, but the source counts are clearly overpredicted. On the basis of these results we propose that clusters of cold gas clouds comprise the majority population of the observed UID EGRET sources. Our interpretation implies that there should be microwave counterparts to most of the UID sources, and will thus be strongly constrained by data from the present generation of microwave anisotropy experiments. ",
        "introduction": "The nature of the UnIDentified (UID) EGRET sources (Hartman et al 1999) is one of the most interesting puzzles in contemporary high energy astrophysics. This puzzle persists primarily because of the dimensions -- typically of order a degree -- of the gamma-ray error boxes: large uncertainties in source position prohibit deep searches for counterparts at other wavelengths. We know, however, that this is only part of the problem: in some cases the EGRET source is bright enough that it can be located quite accurately, yet still no clear counterpart can be found (e.g. Mirabel et al 2000; Reimer et al 2001).  The majority of the discrete sources detected by the EGRET instrument  are UID sources (Hartman et al 1999); it is not clear at present whether they are fundamentally similar to the known GeV emitters, or whether they represent an entirely new population. Because the UID sources are predominantly located at low to mid Galactic latitudes, they must be a Galactic population. This rules out one of the major classes of known discrete celestial GeV gamma-ray sources, namely the flat-spectrum radio quasars as these are extragalactic and hence isotropically distributed on the sky. Much attention has therefore focussed on the possibility that the UID sources are related to massive stars and their evolutionary end-points. In particular there has been intense interest in young pulsars as a contributor to the UID population  (e.g. Bailes and Kniffen 1992; Kaaret and Cottam 1996; Roberts, Romani and Johnston 2001; Grenier and Perrot 2001); several of the identified GeV sources are young pulsars (Hartmann et al 1999). However, the observed latitude distribution of the UID sources is too broad to be readily compatible with young pulsars (Yadigaroglu and Romani 1997). Supernova remnants (Sturner, Dermer and Mattox 1996; Torres et al 2003), molecular clouds and massive star clusters (Montmerle 1979; Benaglia et al 2001) share this difficulty; but Gehrels et al (2001) have noted that the sky distribution of UID sources is consistent with a strong contribution from objects in Gould's Belt, thus supporting a connection with massive stars.  Given the difficulty of detecting counterparts to the UID EGRET population, it is natural to consider the possibility that these sources may be connected with the dark matter problem. The currently popular cosmological model, the Cold Dark Matter model (e.g. Peebles 1993), stipulates that the dark matter is composed of weakly interacting particles; it is possible that these particles are massive -- i.e. they are WIMPS -- and could yield gamma-rays by annihilation or decay. This possibility has been explored by many authors, e.g. Calc\\'aneo-Rold\\'an and Moore (2000), Bergstr\\\"om, Edsj\\\"o and Gunnarsson (2001).  There are of course other types of dark matter. Of particular interest in the present context is baryonic dark matter in the form of cold, dense gas clouds: unlike diffuse gas, cold, dense clouds are difficult to detect directly (Pfenniger, Combes and Martinet 1994; Gerhard and Silk 1996; Combes and Pfenniger 1997; Walker and Wardle 1999). There are many astrophysical motivations for considering dark matter in this form, ranging from the properties of late-type galaxies (Pfenniger, Combes and Martinet 1994; Walker 1999), to observations of radio-wave scintillation (Henriksen and Widrow 1995; Walker and Wardle 1998), and the ``Blank Field'' SCUBA sources (Lawrence 2001). Indirect constraints -- from the small amplitude of fluctuations in the Cosmic Microwave Background (CMB), and from Big Bang Nucleosynthesis  -- are generally supposed to exclude any significant amount of baryonic dark matter (e.g. Turner and Tyson 1999), but these arguments are not entirely free of loopholes (Hogan 1993; Walker and Wardle 1999) and direct constraints are desirable. There is not yet any direct observational evidence which excludes a substantial fraction of the Galactic dark matter being in the form of cold, dense gas clouds (Pfenniger, Combes and Martinet 1994; Gerhard and Silk 1996; Walker and Wardle 1999; Rafikov and Draine 2000; Ohishi, Mori and Walker 2003; Walker and Lewis 2003). But the key motivation here for considering this form of dark matter is that gas clouds are expected to emit gamma-rays if they reside in the cosmic-ray disk (e.g. Bloemen 1989), and any unexplained gamma-ray flux might therefore be a signature of dark matter in this form (De~Paolis et al 1995; Kalberla, Shchekinov and Dettmar 1999; Sciama 2000). The Galactic gamma-ray halo reported by Dixon et al (1998) has, for example, been interpreted in these terms (De~Paolis et al 1999). The low mass ($\\sim10^{-4}\\,{\\rm M_\\odot}$) of the individual clouds renders them undetectable by EGRET when they exist in isolation (see \\S5). However,  clustering is ubiquitous in gravitating systems, and we therefore consider the possibility that the UID EGRET sources are simply clusters of cold, dense gas clouds.  The present paper is organised as follows: we begin by describing the model we have employed for the distribution of cold gas clouds (\\S2), and the emission which is expected from the gas as a result of cosmic-ray interactions (\\S3); our model cosmic-ray distribution is given in \\S4; we present the results of our Monte Carlo simulation in \\S5, and then compare these results to the EGRET data in \\S6. ",
        "conclusions": "If our Galaxy contains a significant component of dark matter in the form of cold, dense gas clouds, clustered into large aggregates, some of those clusters should have been detected by EGRET. Using a CDM-like mass spectrum for the clustering, we have shown that the predicted gamma-ray source population has properties which are broadly similar to those of a large fraction of the UID EGRET sources. In particular: the Galactic latitude distribution and the source size distribution anticipated in the model find support in the EGRET data. Furthermore the intrinsically ``dark'' nature of the predicted sources naturally explains most of the difficulty in finding counterparts; however, the large angular size of the clusters also plays a role, because counterpart searches to date have had poor sensitivity to $\\sim\\,$degree-sized sources. The total number of UID sources predicted by the model is too large, particularly bearing in mind that many of the observed UID sources are likely to be examples of known types of gamma-ray emitters. Data returned by the various CMB experiments will provide a powerful test of the interpretation we have presented: counterpart thermal microwaves should be detected from the cold gas, and the source structure should be resolved."
    },
    "0212/astro-ph0212091_arXiv.txt": {
        "abstract": "We present phase-referenced VLBI observations of the radio continuum emission from, and the neutral hydrogen 21~cm absorption toward, the Ultra-Luminous Infrared Galaxy IRAS~17208--0014. The observations were carried out at 1362~MHz using the Very Long Baseline Array, including the phased Very Large Array as an element. The high-resolution radio continuum images reveal a nuclear starburst region in this galaxy, which is composed of diffuse emission approximately $670 \\times 340$~pc on the plane of the sky, and a number of compact sources. These sources are most likely to be clustered supernova remnants and/or luminous radio supernovae. Their brightness temperatures range over (2.2--6.6)$\\times 10^{5}$~K, with radio spectral luminosities between $(1-10) \\times 10^{21}~{\\rm W~Hz}^{-1}$. The total VLBI flux density of the starburst region is $\\sim$52~mJy, which is about 50\\% of the total flux density detected with the VLA at arcsecond resolution. For this galaxy, we derive a massive star formation rate of $\\sim$84$\\pm 13~M{_\\odot}$~yr$^{-1}$, and a supernova rate of $\\sim$4$\\pm$1~yr$^{-1}$. H~{\\footnotesize I} absorption is detected in multiple components with optical depths ranging between 0.3 and 2.5, and velocity widths between 58 and 232~km~s$^{-1}$. The derived column densities, assuming $T_{\\rm s}=100~\\rm {K}$, range over $(10-26) \\times 10^{21}$~cm$^{-2}$. The H~{\\footnotesize I} absorption shows a strong velocity gradient of 453~km~s$^{-1}$ across $0\\rlap{.}''36$ (274~pc). Assuming Keplerian motion, the enclosed dynamical mass is about {$2.3 \\times 10^9~({\\rm sin}^{-2}i)~M{_\\odot}$}, comparable to the enclosed dynamical mass estimated from CO observations. ",
        "introduction": "At luminosities above $10^{11}~L{_\\odot}$, infrared galaxies become the most numerous objects in the local universe ($z \\leq 0.3$) \\citep {SM96}. The trigger for the intense infrared emission appears to be the strong interaction or merger of molecular gas-rich spirals. Galaxies at the highest infrared luminosities ($L{_{\\rm IR}}(8-1000~{\\mu}m) \\geq 10^{12} L{_\\odot}$), known as Ultra-Luminous Infrared Galaxies (ULIRGs), appear to be advanced merger systems and may represent an important stage in the formation of quasi-stellar objects \\citep{SAN88}.  The bulk of the energy radiated by these sources is infrared emission from warm dust grains heated by a central power source or sources. The critical question concerning these galaxies is whether the dust is heated by a nuclear starburst, or an active galactic nucleus (AGN), or a combination of both. Mid-infrared spectroscopic studies on a sample of ULIRGs by \\citet {GEN98}, suggest that 70\\%--80\\% of these galaxies are powered predominantly by recently formed massive stars, and 20\\%--30\\% by a central AGN. These authors conclude further that at least half of these ULIRGs are probably powered by both an AGN and a starburst in a 1--2~kpc diameter circumnuclear disk or ring. The most direct evidence to date of sub-kpc nuclear starburst regions in ULIRGs is the discovery of luminous radio supernovae and supernova remnants in both Arp~220 \\citep {SMI98} and Mrk~273 \\citep {CT00}, using VLBI observations.  In this paper, we present VLBI continuum and H~{\\footnotesize I} absorption observations on the ULIRG IRAS~17208--0014 at $z=0.0426$. This galaxy has an infrared luminosity of $L_{\\rm IR}= 2.5 \\times 10^{12} L{_\\odot}$, as defined in \\citet {GOL95}.  Optical images of IRAS~17208--0014 at 6550 {\\AA} show two tidal tails from a merger \\citep {MM90,MUR96}, and its optical spectrum resembles that of an H~{\\footnotesize II} region \\citep {VEI95,SOI00}. Its near-IR images show a very disturbed morphology and an extended but single nucleus, suggesting a complete merger \\citep {ZL93,MUR96}. Higher resolution near-IR images reveal numerous extremely luminous clusters in the inner 1~kpc, with  no direct evidence of an AGN \\citep {SCO00}. Infrared observations, in general, suggest that this galaxy represents the extreme of starburst dominated sources of this type \\citep {SOI00}, and an observational proof that a collision of galaxies can lead to a mass distribution similar to elliptical galaxies \\citep {ZL93}.  The CO (1--0) emission observations of IRAS~17208--0014 by \\citet {DS98} show a source of size $1\\rlap{.}''8 \\times 1\\rlap{.}''6$ at full width half maximum, and a strong velocity gradient with a change of 400~km~s$^{-1}$ over $1\\rlap{.}''5$ in position angle $120^{\\circ}$. The enclosed dynamical mass estimated from the CO observation is $8.5 \\times 10^{9}~M{_\\odot}$.  IRAS~17208--0014 also exhibits OH megamaser activity \\citep {MAR89} at 1665, 1667, and 1720~MHz. The strongest emission is in the 1667~MHz line, with a luminosity of $L_{\\rm OH}= 10^{3}~L{_\\odot}$. MERLIN 1662~MHz radio continuum observations at $0\\rlap{.}''3$ resolution \\citep {MAR89} revealed the existence of an unresolved component on the longest baselines, with a constant visibility amplitude of $\\sim$35~mJy. However, the continuum emission from this galaxy was not detected in the 18~cm global VLBI observations by \\citet {DLLS99} to a limit of $3\\sigma=210~\\mu{\\rm Jy~beam}^{-1}$.  Single dish observations by \\citet {MAR89} at Nan\\c{c}ay showed the existence of a very wide H~{\\footnotesize I} absorption line in this galaxy, with full velocity widths of 650 and 695~km~s$^{-1}$ at 50\\% and 20\\% of the maximum depth, which is $-12$~mJy.  In this paper, we report a detailed study of the H~{\\footnotesize I} 21~cm absorption and the radio continuum emission from IRAS~17208--0014. The results reveal the H~{\\footnotesize I} gas dynamics and the distribution of compact continuum sources, possibly composed of luminous radio supernova and/or supernova remnants, in the nuclear region of this galaxy on sub-kpc scales. We adopt a distance of 171~Mpc to this galaxy, assuming ${H_\\circ=75}$~km~s$^{-1}$~Mpc$^{-1}$. At this distance 1~mas corresponds to 0.76~pc. ",
        "conclusions": "We have presented the results of phase-referenced VLBI observations, using the VLBA and the phased VLA, of the 21~cm continuum emission and the H~{\\footnotesize I} absorption in the central $\\sim$1.1~kpc of the very advanced merger galaxy IRAS~17208--0014.  The high resolution continuum images reveal the details of the previously undetected nuclear starburst region of this galaxy. Both diffuse and compact continuum emission are detected. The total VLBI flux is 52~mJy, and represents only half of the total flux seen with the VLA at a lower resolution, suggesting the existence of diffuse emission not detected by our VLBI array. The compact sources in the starburst region are more likely clustered luminous radio supernovae and supernova remnants, considering the H~{\\footnotesize II} spectra observed in this galaxy and its morphological similarities to other well-known starbursts. However, we cannot rule out the possibility that each of our compact sources is mainly powered by an individual bright RSN nested in a region that contains faded SNRs. The brightness temperatures of the compact structures seen in these observations are $>10^5$~K. The flux density of the brightest source in the starburst region is less than 3\\% of the total radio flux density. The results suggest that there is no radio-loud AGN in the nuclear region of IRAS~17208--0014.  Both the far-infrared luminosity and the radio continuum flux of this galaxy imply a massive star formation rate of $\\sim$$84~M_\\odot~{\\rm yr}^{-1}$ and supernova rate of $\\sim$$4~{\\rm yr}^{-1}$. The estimated individual luminous RSN number agrees surprisingly well with the number of the compact sources detected above 5$\\sigma$ level. From the minimum energy condition, we have estimated the pressure to be $\\sim$$6 \\times 10^{-10}$~dyn~cm$^{-2}$ and the magnetic field to be $\\sim$$144~\\mu$Gauss in the starburst region.  The very wide H~{\\footnotesize I} absorption is composed of five components with velocity widths between 58 and 232~km~s$^{-1}$. The column densities of these absorption peaks are on the order of $10^{20}~T_{\\rm s}{\\rm (K)}$, and the derived visual extinctions are between $0.066~T_{\\rm s}{\\rm (K)}$ and $0.174~T_{\\rm s}{\\rm (K)}$~mag. The strongest velocity gradient in the H~{\\footnotesize I} absorption is seen in position angle 120$^{\\circ}$, as in the CO $(1-0)$ emission. The H~{\\footnotesize I} position-velocity profile in this position angle shows a velocity gradient of 1653~km~s$^{-1}$~kpc$^{-1}$. The calculated dynamical mass, assuming Keplerian motion, is {$2.3 \\times 10^9~({\\rm sin}^{-2}i)~M{_\\odot}$}, comparable to the enclosed mass obtained from CO observations. The H~{\\footnotesize I} absorption results suggest the existence of a neutral gas disk containing several clouds, situated inside a larger scale molecular disk as seen in the interferometric CO $(1-0)$ emission line observations."
    },
    "0212/cond-mat0212223_arXiv.txt": {
        "abstract": "The formation of large-scale vortices is an intriguing phenomenon in two-dimensional turbulence. Such organization is observed in large-scale oceanic or atmospheric flows, and can be reproduced in laboratory experiments and numerical simulations. A general explanation of this organization was first proposed by Onsager (1949) by considering the statistical mechanics for a set of point vortices in two-dimensional hydrodynamics. Similarly, the structure and the organization of stellar systems (globular clusters, elliptical galaxies,...) in astrophysics can be understood by developing a statistical mechanics for a system of particles in gravitational interaction as initiated by Chandrasekhar (1942). These statistical mechanics turn out to be relatively similar and present the same difficulties due to the unshielded long-range nature of the interaction. This analogy concerns not only the equilibrium states, i.e. the formation of large-scale structures, but also the relaxation towards equilibrium and the statistics of fluctuations. We will discuss these analogies in detail and also point out the specificities of each system. ",
        "introduction": "\\label{sec_introduction}   Two-dimensional flows with high Reynolds numbers have the striking property of organizing spontaneously into \\index{Coherent structures} coherent structures (the vortices) which dominate the dynamics \\cite{williams} (see Fig. \\ref{McW}). The robustness of Jupiter's Great Red Spot, a huge vortex persisting for more than three centuries in a turbulent shear between two zonal jets, is probably related to this general property. Some other coherent structures like dipoles (pairs of cyclone/anticyclone) and sometimes tripoles have been found in atmospheric or oceanic systems and can persist during several days or weeks responsible for atmospheric blocking. Some astrophysicists invoke the existence of organized vortices in the gaseous component of disk galaxies in relation with the emission of spiral density waves \\cite{nezlin}. It has also been proposed that planetary formation might have begun inside persistent gaseous vortices born out of the protoplanetary nebula \\cite{barge,tanga,bracco,godon,cplanetes} (see Fig. \\ref{planetes}). As a result, hydrodynamical vortices occur in a wide variety of geophysical or astrophysical situations and their robustness demands a general understanding.  Similarly, it is striking to observe that self-gravitating systems \\index{Astrophysics} follow a kind of organization despite the diversity of their initial conditions and their environement \\cite{bt} (see Fig. \\ref{Hubble}). This organization is illustrated by morphological classification schemes such as the Hubble sequence for galaxies and by simple rules which govern the structure of individual self-gravitating systems. For example, elliptical galaxies display a quasi-universal luminosity profile described by de Vaucouleur's $R^{1/4}$ law and most of globular clusters are well-fitted by the Michie-King model. On the other hand, the flat rotation curves of spiral galaxies can be explained by the presence of a dark matter halo with a density profile decreasing as $r^{-2}$ at large distances. The fractal nature of the interstellar medium and the large scale structures of the universe also display some form of organization.   \\begin{figure} \\centering \\includegraphics[angle=0,width=0.8\\textwidth]{chava-fig1.eps} \\caption[]{Self-organization of two-dimensional turbulent flows into large-scale vortices \\cite{williams}. These vortices are long-lived and dominate the dynamics.} \\label{McW} \\end{figure}  \\begin{figure} \\centering \\includegraphics[angle=0,width=0.9\\textwidth]{super.eps} \\caption[]{A scenario of planet formation inside large-scale vortices presumably present in the Keplerian gaseous disk surrounding a star at its birth. Starting from a random vorticity field, a series of anticyclonic vortices appears spontaneously (upper panel). Due to the Coriolis force and to the friction with the gas, these vortices can efficiently trap dust particles passing nearby (lower pannel). The local increase of dust concentration inside the vortices can initiate the formation of planetesimals and planets by gravitational instability. This numerical simulation is taken from \\cite{bracco}.} \\label{planetes} \\end{figure}    \\begin{figure} \\centering \\includegraphics[angle=0,width=0.55\\textwidth]{superG.ps} \\caption[]{Large-scale structures in the universe as observed with the Hubble space telescope. The analogy with Fig. \\ref{McW} is striking and will be discussed in detail in this paper.} \\label{Hubble} \\end{figure}  The question that naturally emerges is what determines the particular configuration to which a self-gravitating system or a large-scale vortex settles. It is possible that their actual configuration crucially depends on the conditions that prevail at their birth and on the details of their evolution. However, in view of their apparent regularity, it is tempting to investigate whether their organization can be favoured by some fundamental physical principles like those of thermodynamics and statistical physics. We ask therefore if the actual states of self-gravitating systems in the universe and coherent vortices in two-dimensional turbulent flows are not simply more probable than any other possible configuration, i.e. if they cannot be considered as {\\it maximum entropy states}. This statistical mechanics approach has been initiated by Onsager \\cite{onsager} for a system of point vortices and by Chandrasekhar \\cite{chandra42} in the case of self-gravitating systems.  It turns out that the statistical mechanics of two-dimensional vortices and self-gravitating systems present a deep analogy despite the very different physical nature of these systems. This analogy was pointed out by Chavanis in \\cite{cthese,cfloride,japon} and further developed in \\cite{csr,drift,kin1,kin2,csire1,csire2}. In the following, we will essentially discuss the statistical mechanics of 2D vortices and refer to the review of Padmanabhan \\cite{pad} (and his contribution in this book) for more details about the statistical mechanics of self-gravitating systems. We will see that the analogy between two-dimensional vortices and (three-dimensional) self-gravitating systems concerns not only the prediction of the equilibrium state, i.e. the formation of large-scale structures, but also the statistics of fluctuations and the relaxation towards equilibrium.  This paper is organized as follows. In Sec. \\ref{sec_statmech}, we discuss the statistical mechanics of point vortices introduced by Onsager \\cite{onsager} and further developed by Joyce \\& Montgomery \\cite{jm} and Pointin \\& Lundgren \\cite{pl} among others (see a complete list of references in the book of Newton \\cite{newton}). We discuss the existence of a thermodynamic limit in Sec. \\ref{sec_field} and make the connexion with field theory. Statistical equilibrium states of axisymmetric flows are obtained analytically in Sec. \\ref{sec_disk}-\\ref{sec_unbounded}. The relation with equilibrium states of self-gravitating systems is shown in Sec. \\ref{sec_gravN}. In Sec. \\ref{sec_fluc}, we discuss the statistics of velocity fluctuations produced by a random distribution of point vortices and use this stochastic approach to obtain an estimate of the diffusion coefficient of point vortices. Application to 2D decaying turbulence is considered in Sec. \\ref{sec_anomalous}. In Sec. \\ref{sec_bath1}, we describe the relaxation of a point vortex in a thermal bath and analyze this relaxation in terms of a Fokker-Planck equation involving a diffusion and a drift. In Sec. \\ref{sec_kin}, we develop a more general kinetic theory of point vortices. A new kinetic equation is obtained which satisfies all conservation laws of the point vortex system and increases the Boltzmann entropy (H-theorem). We mention the connexion with the kinetic theory of stars developed by Chandrasekhar \\cite{chandra42}. In Sec. \\ref{sec_violrelax}, we discuss the violent relaxation of 2D vortices and stellar systems. We mention the analogy between the Vlasov and the Euler equations and between the statistical approach developed by Lynden-Bell \\cite{lb} for collisionless stellar systems and by Kuz'min \\cite{kuzmin}, Miller \\cite{miller} and Robert \\& Sommeria \\cite{rs1} for continuous vorticity fields. The concepts of ``chaotic mixing'' and ``incomplete relaxation'' are discussed in the light of a relaxation theory in Sec. \\ref{sec_MEPP}. Application of statistical mechanics to geophysical flows and Jupiter's Great Red Spot are evocated in Sec. \\ref{sec_del}. ",
        "conclusions": "\\label{sec_conclusion}   The statistical mechanics of 2D vortices and \\index{Astrophysics} stellar systems appear to be remarkably similar despite the different nature of these systems. We have tried to develop this analogy in different directions. First of all, the $N$-star and $N$-vortex problems both involve an unshielded long-range potential generated by the density of particles themselves. At statistical equilibrium, these systems are described by a Boltzmann-Poisson equation whose solutions characterize organized states (at negative temperatures for vortices). In the case of stars, the relaxation towards equilibrium can be viewed as a competition between a diffusion and a friction. We have proposed to describe the relaxation of point vortices similarly in terms of a diffusion and a drift. The diffusion is due to the fluctuations of the force experienced by a star or to the fluctuations of the velocity field moving a vortex. The statistics of these fluctuations can be studied by similar mathematical methods. The fluctuations of the gravitational field are described by a particular L\\'evy law, called the Holtzmark distribution, and the fluctuations of the velocity of vortices are described by a marginal Gaussian distribution, intermediate between Gaussian and L\\'evy laws. The friction experienced by a star is due fundamentally to the inhomogeneity of the velocity distribution of the stellar cloud. Analogously, the drift experienced by a vortex results from the spatial inhomogeneity of the vortex cloud. The friction and the drift can be understood similarly in terms of a polarization process and a back reaction of the system. In the thermal bath approximation, the coefficients of friction and drift are given by an Einstein relation and the one-body distribution function satisfies a Fokker-Planck equation. Further away from equilibrium, the collisional dynamics of stars is described by the gravitational Landau equation. We have derived a new kinetic equation that should be appropriate to the ``collisional'' relaxation of point vortices.  A system of stars or vortices can also undergo a form of violent relaxation. This is essentially a collisionless process driven by the rapid fluctuations of the potential as a result of collective effects (chaotic mixing). In this collisionless regime, the stars and the vortices are described by the Vlasov-Poisson and the Euler-Poisson systems. These equations present a similar structure and a statistical mechanics can be developed to predict the ``most probable state'' resulting from a complex evolution driven by a mixing process. The relaxation towards equilibrium can be incomplete and an out-of-equilibrium study is necessary to understand what limits relaxation and causes a kinetic confinement of the system in a ``maximum entropy bubble''. Relaxation equations have been derived either from a heuristic Maximum Entropy Production Principle or from a more controllable kinetic theory, in an asymptotic regime of the dynamics (gentle relaxation) in which a quasilinear approximation can be implemented.  It should be noted that many results presented in this paper, in particular those corresponding to the kinetic theory of 2D vortices developed in Secs. \\ref{sec_fluc}-\\ref{sec_kin}, are very recent and need to be completed and further discussed. In particular, it appears indispensable to carry out extensive numerical simulations to test their relevance and determine their domains of applicability. It is plausible that the true dynamics of stars and vortices is more complex than the picture that has been given here. In addition, the description of chaos in these systems has not been addressed at all in this paper although it is presumably an essential ingredient to understand their dynamics. We are thus far from reaching a complete understanding of these systems with long-range interactions. We feel, however, that the analogy between the statistical mechanics of stars and vortices that we have investigated is correct in its mains lines and should lead again to fruitful developements."
    },
    "0212/astro-ph0212022_arXiv.txt": {
        "abstract": "The energy range of hard X-rays is a key waveband to the study of high energy processes in celestial objects, but still remains poorly explored. In contrast to direct imaging methods used in the low energy X-ray and high energy gamma-ray bands, currently imaging in the hard X-ray band is mainly achieved through various modulation techniques. A new inversion technique, the direct demodulation method, has been developed since early 90s. With this technique, wide field and high resolution images can be derived from scanning data of a simple collimated detector. The feasibility of this technique has been confirmed by experiment, balloon-borne observation and analyzing simulated and real astronomical data. Based the development of methodology and instrumentation, a high energy astrophysics mission -- Hard X-ray Modulation Telescope (HXMT) has been proposed and selected in China for a four-year Phase-A study. The main scientific objectives are a full-sky hard X-ray (20-200 keV) imaging survey and high signal-to-noise ratio timing studies of high energy sources. ",
        "introduction": "The hard X-ray band  with photon energies from $\\sim 10$~keV to a few hundred keV is very important for the study of high energy processes near compact objects, super-massive black holes and in relativistic outflows. Astronomical imaging in hard X-rays is still an observational problem. This is because high energy photons in this range can neither be focused like lower energy photons, nor can their arrival direction be determined from Compton scattering or e$^{\\pm}$ pair production as in the high energy $\\gamma$-ray range. Due to the technological difficulties no detailed all-sky map of hard X-rays has been provided yet.  In contrast to direct imaging methods used in the optical and soft X-ray bands, currently imaging in the hard X-ray band is mainly achieved through various modulation techniques. M. Oda (1965) proposed that a collimator composed of two parallel grids of wires, scanning cross a X-ray source over a detector, could reveal the positional information of the source. After then the rotation modulation collimator (RMC) technique (Schnopper et al. 1968; Willmore 1970) and coded aperture mask (CAM) technique (Dick 1968; Ables 1968) were developed with good angular resolution and a wide field of view (FOV) for hard X-ray imaging. By rotating the modulation grids, a periodic modulation of the source flux is introduced which contains unique information on the position and intensity of the sources in the FOV. Images are obtained with the aid of Fourier analysis or cross-correlation analysis. Composite-aperture telescopes, e.g. multipitch telescopes, may be used in order to reduce the side lobes and perform imaging of diffuse sources. Currently, in hard X- ray imaging observations, the CAM is widely used. Modulated spatially by the coded aperture mask, the incident photons are recorded by a position sensitive detector. The true spatial intensity distribution is reconstructed from the observational data by using the cross-correlation analysis method or other mathematical decording techniques. The coded aperture mask technique is most widely used in hard X-ray imaging: modulated spatially by the coded aperture mask, the incident photons are recorded by a position sensitive detector. The true objective distribution is reconstructed from the observational data by using the cross-correlation (CC) technique or other inversion methods. Due to the position sensitive detector a hard X-ray coded aperture telescope is fairly complicated and expensive. The intrinsic angular resolution of a coded aperture mask instrument can be estimated by $\\Delta = \\arctan (2\\delta x/d)$, where $\\delta x$ is the detector-mask pixel size and $d$ the distance between the mask and detector plane. With a limited position resolution of the detector, say 1 cm, the telescope can reach a length of $6 \\sim 7$ m to get about $10'$ angular resolution, which leads to large structure, heavy weight and difficulties in suppressing the local background. The effective area is reduced to half or more due to the occultation of the mask and a wide field of view will lead to a higher background through the aperture. Complicated image distortions (e.g. side lobes, pseudo images etc.) are usually appeared in images obtained by coded aperture telescopes, which strongly reduce the angular resolution and imaging capability for weak sources.  Since early 1990s a direct demodulation method has been developed. The direct demodulation technique can extract the information on the observed object from the observational data more sufficiently than conventional inversion techniques, with which one can perform wide field and high resolution imaging with a non-position-sensitive detector and relatively simple modulation methods. ",
        "conclusions": "Sufficiently using the information on the observed object is of great importance in dealing with the reconstruction problem. The observation equation system \\[\\hspace{30mm} \\sum\\limits_{i=1}^{N}p(k,i)f(i)=d(k) \\hspace{10mm} (k=1,...,M) \\hspace{31mm}(1) \\] or \\[\\hspace{60mm} Pf=d \\hspace{56mm}(2) \\] contains all information about the object $f$ in the observed data $d$ obtained by the observation with an instrument $P$. The maximum-entropy method (ME) is a widely used method of reconstructing object $f$ from observed data $d$. ME is to choose a solution $f(i) (i=1,...,N)$ from all satisfying the statistical criterion  \\begin{equation} \\sum_{k=1}^M\\{[\\sum_ip(k,i)f(i)-d(k)]/\\sigma_k\\}^2=M \\end{equation} by the condition that the information entropy \\begin{equation} -\\sum_if(i)\\log f(i)=\\max \\end{equation} Many informations on the object $f$ are lost when using the single equation (9) instead of the equation system (1) containing $M$ equations, that infinite intensity distribution $f$ can satisfy Eg. (9). The maximum-entropy condition (10) helps us to pick out a most smooth one from them, but not necessary the best one. Another widely used method is the cross-correlation technique, which takes the cross-correlation distribution $c$ of the data and the modulation pattern to estimate the object intensity \\begin{equation} f\\propto c=P^Td \\end{equation} From the correlation equations \\[ P'f=c   (5) \\] one can see $c\\ne f$, the cross-correlation $c$ is just an image of the object $f$ through a modulation (distortion) of a certain imaging instrument with a point spread function $P'$. The cross-correlation reconstruction completely ignores the information included in the correlation equations, still does not make full use of the information on the object obtained by the observation. In general the two kinds of inversion methods have no significant difference in their results, the resultant resolution are limited by the intrinsic resolution $\\Delta_0$ of the instrument \\begin{equation} \\Delta_{CC}\\approx \\Delta_{ME}\\approx \\Delta_0 \\end{equation}  The direct decording approach, reconstructing objects by iteratively solving observational equations or correlation equations under nonlinear constraints, each iterative calculation is directly based on the modulation equation. Obviously the direct approach can use the information containing in the equations more sufficiently than any indirect one through a statistical criteria or cross-correlation transformation and then can obtain much better resolution $\\Delta_{DD}\\ll \\Delta_0$.  With conventional techniques the resolution is limited by the intrinsic one, it does not matter how sensitive the observation is.  For improving the precision and resolution of an experiment, people are usually devoted to improving the intrinsic resolution by constructing high-precision instrument, more and more complicated and high-priced. It is a most difficult task to conduct an experiment with both high resolution and high sensitivity. Under the direct inversion technique, the resultant resolution $\\Delta_{DD}$ is not only dependent on the instrumental intrinsic resolution $\\Delta_0$, but also dependent on the amount $S$ of observed signal, the ratio of signal to noise $SNR$ and modulation pattern $M$ \\begin{equation} \\Delta_{DD}=f(\\Delta_0, S, SNR, M) \\end{equation} Applying the direct demodulation technique to space instrument design, we can realize high resolution imaging with a low resolution or even non-position-sensitive detector, and can conduct observations with both high resolution and high sensitivity by improving the resolution ability through increasing the detection sensitivity, as shown by the expected performance of HXMT mission.  The main difficult of the direct methods is the illness of solution causing by errors in data and other uncertainties in modulation equations. Introducing physical constraints to make nonlinear control in the iterative process of solving modulation equations is a key point to make the direct demodulation realizable, and, furthermore, stable, convergent and global optimal. Benefited from the physical constraints in iterates the effect of errors can be depressed effectively in the demodulation process. As a consequence, the direct demodulation can get images with much less false appearances and be capable of reconstructing not only discrete sources, but also extended features as well."
    },
    "0212/astro-ph0212508_arXiv.txt": {
        "abstract": "2$\\mu$ broad- and narrow-band imaging with the Gemini-N telescope of five z$\\sim$4.7 QSOs, has resolved both the host galaxies and [O II] emission-line gas. The resolved fluxes of the host galaxies fall within the extrapolated spread of the K-z relationship for radio galaxies at lower redshifts, and their resolved morphology is irregular. The [O II] images indicate knots coincident with many continuum features and also some bright jet-like features near the nucleus. The line emission total fluxes indicate overall equivalent widths of 5 to 10\\AA~ at rest wavelengths. Two of the QSOs are in a local environment of faint galaxies of similar magnitude to the hosts, and three have nearby galaxies with excess narrow-band flux, which would be [O~II] if they are at the QSO redshift. ",
        "introduction": "The host galaxies of high redshift QSOs are of interest as they are considered to be massive galaxies with high abundances and massive central black holes in the early universe. It has been found that at redshift 2 and higher, QSO host galaxies are large, luminous, and the sites of star-formation and galaxy merging. There have been several investigations of QSO host galaxies in the redshift range 2 to 2.5 (e.g. Ridgway et al 2001, Kukula et al 2001, Hutchings et al 2002). There have been studies of radio galaxies, which are easier to resolve, to redshifts 4 and higher (e.g.  van Breugel et al 1998, Inskip et al 2001, Jarvis et al 2001, De Breuck et al 2002). Compared with radio galaxies, QSO hosts are more difficult to study because of the bright nuclear light source, but recent surveys, in particular the SDSS, have greatly increased the number of QSOs as known high redshift objects. This paper describes observations intended to detect and characterize some QSO hosts at redshifts near 4.7.   Most of the work on higher redshift QSOs has been done with HST or ground-based AO systems, to reduce the light scattered by the QSO nucleus. While these studies have found that there are compact spheroidal stellar populations around the nuclei (see above references), there are also indications that the host galaxies have faint light extended over several arcsec (e.g. Hutchings 1995). Assuming that the hosts are still large and luminous at redshifts 4.7, the main requirement was of large light-gathering power, since redshift dimming will be large. Subarcsecond image quality is still very desirable, to avoid too much contamination by the bright unresolved nucleus. Finally, in order to learn about both stars and gas in the galaxies, the observations were made in broad-band and narrow-band at the redshifted wavelength of the [O II] emission line, characteristic of star-formation. Thus, the rest wavelengths are near 3700\\AA, and the observations made near 2 microns. Similar narrow-band imaging of the redshifted [O III] lines is described by Hutchings, Morris, and Crampton (2001).  The QSOs were selected from a high redshift list of Djorgovski (http://www.astro.caltech.edu/$\\sim$george/z4.qsos) based mainly on the SDSS (e.g. Anderson et al 2001, Zheng et al 2000). The sample selected has redshifts that place the [O~II] line inside the bandpass of a NIRI narrow-band filter. Objects were selected from this short list by the Gemini service observers, to fit their schedules and weather conditions. Thus, the sample is essentially random in all properties other than redshift being close to 4.7, and the magnitude limits of the SDSS.   The observations were carried out on the Gemini N telescope during the first half of 2002, with total exposures of 40 minutes in both broad- and narrow-band filters. The details are given in Table 1, along with properties of the QSOs, including the R/r magnitude and QSO names from Djorgovski's list, and some measured quantities. As the observations were carried out by service observing over a number of nights, the conditions varied somewhat between them. The FWHM and limiting surface brightness are shown for all objects, and lie in reasonably compact ranges of values. We refer below to the narrow- and broad-band observations as NB and BB respectively.  The analysed images were derived from dithered exposures, processed using standard techniques. Details of this process for other ground-based observations are given in Hutchings et al (1999), and Steinbring, Crampton, and Hutchings (2002). The field size is approximately 2 arcmin, which included suitable PSF stars in all cases, and enabled some analysis of the faint galaxy counts in the QSO neighbourhood. ",
        "conclusions": ""
    },
    "0212/astro-ph0212214_arXiv.txt": {
        "abstract": "{We reanalyze the strong lens modeling of the cluster of galaxies MS2137.3-2353 using a new $UBVRIJK$ data set obtained with the ESO Very Large Telescope. We infer the photometric redshifts of the two main arc systems which are both found to be at $z=1.6\\pm 0.1$. After subtraction of the central cD star light in the previous F702/HST imaging we found only one object lying underneath. This object has the expected properties of the fifth image associated to the tangential arc. It lies at the right location, shows the right orientation and has the expected signal-to-noise ratio.  We improve the previous lens modelings of the central dark matter distribution of the cluster, using two density profiles: an isothermal model with a core, and the NFW-like model with a cusp. Without the fifth image, the arc properties together with the shear map profile are equally well fit by  the isothermal model and  by a sub-class of generalized-NFW mass profiles having inner slope power index in the range $0.7\\leq \\alpha \\leq 1.2$. Adding new constrains on the center lens position provided by the fifth image favors isothermal profiles that better predict the fifth image properties. A detailed model including nearby cluster galaxy perturbations or the effect of the stellar mass distribution to the total mass inward does not change our conclusions but imposes the $M/L_I$ of the cD stellar component is below 10 at a 99\\% confidence level.  Using our new detailed strong+weak lensing model together with Chandra X-ray data and the cD stellar component we finally discuss intrinsic properties of the gravitational potential. Whereas X-ray and dark matter have a similar orientation and ellipticity at various radius, the cD stellar isophotes are twisted by $13^\\circ \\pm 3^\\circ$. The sub-arc-second  azimuthal shift we observe between the radial arc position and the predictions of elliptical models correspond to what is expected from a mass distribution twist. This shift may result from a projection effect of the cD and the cluster halos, thus revealing the triaxiality of the mass components. ",
        "introduction": "Cosmological N-body simulations of hierarchical structures formation in a universe dominated by collision-less dark matter predict universal density profiles of halos that can be approximated by the following distribution \\be{nfw} \\rho(r)=\\rho_s \\left(\\tfrac{r}{r_s}\\right)^{-\\alpha} \\left[1+\\tfrac{r}{r_s}\\right]^{\\alpha-3}. \\ee The early simulations of \\citep{NFW} (hereafter NFW) found $\\alpha = 1$, leading to profiles with a central {\\sl cusp} $\\alpha$ and an asymptotic $r^{-3}$ slope, steeper than isothermal (hereafter IS).  More recently, simulations with higher mass resolution confirmed that the density profile Eq. \\eqref{nfw} can fit the dark matter distribution of halos, although different values of $\\alpha$ were obtained by various authors \\citep[see {\\sl e.g.~}][]{Ghigna00,Bullock01}.  While the collision-less $\\Lambda CDM$ cosmology explains observations of the universe on large scales, two issues concerning these halos are still debated. The first one is the apparent excess of sub-halos predicted in numerical simulations, compared to the number of satellites in halos around normal galaxies \\citep{Klypin99, Moore99}. This discrepancy may be resolved if some of the sub-halos never formed stars in the past and are therefore dark structures \\citep{BKW01,Verde02}. \\citet{Metc01,Keea,Keec} or \\citet{Dalal02} argued that we may already see effects of such dark halos through the perturbations they induce on the magnification on the gravitational pairs of distant QSOs.\\\\ The second prediction is the existence of a cuspy universal profile which cannot explain the rotation curves of dwarf galaxies \\citep{Saluccietal00}. If these discrepancies are not simply due to a resolution problem of numerical simulations, then, as it was pointed out by several authors, they may illustrate a small-scale crisis for current CDM models \\citep{NavSt}. In order to solve these issues, alternatives to pure collision-less cold dark matter particles, have been proposed \\citep{Spergel00,bode01}. Also several physical mechanisms which could change the inner slope of mass profiles, like central super-massive black holes \\citep{Merritt,Haehnelt02}, tidal-merging processes inward massive halos \\citep{Dekel02} or adiabatic compression of dark matter can be advocated \\citep[see {\\sl e.g.}~][]{Blumenthal,Keea}.  The demonstration that halos do follow a NFW mass profile over a wide range of mass scale would therefore be a very strong argument in favor of collision-less dark matter particles. Unfortunately, and despite important efforts, there is still no conclusive evidence that observations single out the universal NFW-likes profile and rule out other models. Clusters of galaxies studies are among the most puzzling. In general, weak lensing analysis or X-rays emission models show that both singular isothermal sphere (SIS) and NFW fit equally well their dark matter profile, but there are still contradictory results which seem to rule out either NFW or IS models \\citep[see for example][]{Allen98, Tysonetal98, Mellier99, Cloweetal00, Willick00, Cloweetal01, Arabadjis02, Athreya02}. This degeneracy is explained  because most observations probe the density profile at intermediate radial distances, where an IS and a NFW profiles have a similar $r^{-2}$ behavior.  A promising attempt to address the cusp-core debate is to model gravitational lenses with multiple arcs which are spread at different radial distances, where the SIS and the NFW slopes may differ significantly. As emphasized by \\citet{Miralda95}, ideal configurations are clusters with a simple geometrical structure (no clumps) and with the measurements of the stellar velocity dispersion profile of its central galaxy \\citep[See e.g.][]{Kelson02}. The  MS2137.3-2353 cluster satisfies these requirements and turns out to be an exceptional lensing configuration with several lensed images, including a demagnified one we find out in this work at the very center of the lensing potential. In this paper, we analyze the possibility to break the degeneracy between IS and NFW mass profiles using new data set of MS2137.3-2353 obtained at the VLT and the properties of this new fifth image.  The paper is organized as follows. In Sect. \\ref{observ_section} we review the cluster properties after a summary on previous modelings that claimed for very deep photometric observations. This section also presents the new VLT observations and describes the optical properties of the cluster. Sect. \\ref{modeling_section} presents the strong lensing models for softened IS elliptical halos and NFW cuspy profiles. We discuss the global agreement of both approaches within the CDM paradigm in Sect. \\ref{discus_section}. We stress the importance of the detection of the fifth central demagnified image of the tangential arc system and discuss the observational prospects for the near future in Sect. \\ref{conclu_section}. Throughout this paper, we assume a $\\Omega_0 = 0.3,\\;\\Omega_\\Lambda = 0.7$, and $ H_0 \\equiv 100\\;h\\;{\\rm km}\\:{\\rm s}^{-1}\\:{\\rm kpc}^{-1}$ cosmology in which case $ 1 \\arcsec = 3.24\\hmkpc$ at the cluster redshift $z=0.313$. ",
        "conclusions": "By using strong and weak lensing analysis of HST and new VLT data of MS2137.3-2353 we found important new features on the lensing configuration:  The photometric redshifts or the radial and the tangential arcs are both at $1.6\\pm0.1$ in excellent agreement with the recent spectroscopic observations of \\citeauthor{Sand}. \\\\ The extraction of the cD diffuse stellar light has permitted to detect only one single object which turns out to be at the expected position of the fifth image. Furthermore, its orientation, its ellipticity, its signal-to-noise ratio and its morphology correspond to those expected by the lens modeling. Unfortunately, the poor determination of its shape properties hampers the use of its geometry as a local estimate of the magnification matrix toward the center. \\\\ Using the fifth image together with the weak lensing analysis of VLT data, we then improved significantly the lens modeling. The radial mass profile can then be probed over three orders in distance. This additional constraint seems to favor isothermal profiles with flat core or generalized-NFW profiles with $0.8 \\leq \\alpha \\leq 1.2$ without introducing the fifth image knowledge. When this constraint is added together with the prior motivated that cD center and cluster halo center are the same, we favor flat core softened isothermal spheres. The position of the fifth image is in better agreement with an isothermal model than an NFW mass profile. In addition, it is worth noticing that the kinetics of stars should be analyzed in details, considering a precise distribution function that may depart from the commonly assumed Maxwellian.  We point out a misalignment between the diffuse stellar component major axis and both the lens potential and the X-ray isophotes. We argue it is produced by the triaxial shape of the mass components. This extends the previous demonstration of the ellipticity of the projected dark matter halo. This work is a first attempt to improve strong lensing observables and modeling in order to probe both the central DM cusp/core and the triaxiality of DM halos.  It would be essential to confirm the detection of the fifth image. Fig. \\ref{SED_fig} shows that the spectral energy distribution expected for the fifth image is different than the old stars dominated cD emission. We therefore expect the fifth image to show up on an optimal image subtraction $U-\\lambda J$ ($\\lambda$ being optimized). We attempted to use this technique on our present data but the poor resolution ($\\sim 0.6\\arcsec$) on the U and J ground based images prevent any significant enough detection. We conclude that only a high resolution observation with the Space Telescope in UV-blue wavelengths or in a peculiar emission-line is among the best constraints one could envision in the future.  There is not yet evidence that similar studies as this work can be carried out on other ideal lens configurations. The strength of the diagnostic on the radial mass profile is however so critical that we must apply this technique to a large sample in order to challenge collision-less CDM predictions on a realistic number of clusters of galaxies with eventually a test of the role of dominant central cD galaxies. The simultaneous use of weak lensing data should be more relevant for wider fields in order to check also a $r^{-3}$ fall-off on the density profile predicted at large distance by CDM simulations."
    },
    "0212/astro-ph0212164_arXiv.txt": {
        "abstract": "The properties of radio loud active galaxies (radiogalaxies, BL Lac objects and radio loud quasar) at low and high redshift are briefly reviewed and compared. The recently derived empirical relations between central black hole mass and host properties allow one to explore the demographics of massive black hole in active galaxies. It is shown that all classes of active galaxies considered exhibit very similar host properties and, consequently, have central black holes of comparable mass. The  future observing capabilities would allow the investigation of the evolution of the host galaxy properties up to z $\\sim$ 3, and will give insight into the joint formation of black holes and massive spheroids. ",
        "introduction": "The discovery that many and probably  all apparently inactive nearby early type galaxies harbour a dormant black hole (BH) in their nuclei (e.g. Ferrarese 2002) has changed the view that distinguished active from inactive galaxies. It is also becoming clear that a full comprehension of the processes that produced the galaxies we observe today must take into account the formation of supermassive black holes in their centers.  It is well assessed that the high energy phenomena observed in AGN occur in the nuclei of massive galaxies. Less understood is, however, the link between the properties of the active nucleus and those of the host galaxy. Nevertheless a number of issues are emerging from recent studies. Nuclear activity associated with radio emission is present almost exclusively on galaxies dominated by the spheroidal component. On the other hand radio quiet objects are found in both types of galaxies (bulge and disc dominated galaxies).  A detailed comparison of the host properties of various type of AGNs together with  those of normal (inactive) galaxies can help to investigate the origin of the nuclear activity. In particular one can explore if and how the nuclear activity depends on the global properties (e.g. total luminosity and scale-length) of the galaxies, if there are difference in the stellar population with respect to inactive galaxies and what is the role of interactions (disturbed morphologies, close companions). Such kind of comparison can be done with different degree of details that depends in part on the distance of the sources and on the prominence of the nucleus with respect to the host galaxy. Radiogalaxies are the nearest active objects and exhibit only faint nuclei and are therefore the most easy AGN to study. On the other hand quasars are rare objects locally, have extremely bright nuclei and are rather more difficult to study.  The recent discovery that the BH mass (M$_{BH}$) is correlated with the properties of the bulge component of the host galaxy, which is translated into the relationships between M$_{BH}$ and the bulge luminosity (Kormendy \\& Richstone 1995; Magorrian et al. 1998; Richstone et al. 1998; Kormendy \\& Gebhardt 2001) and between M$_{BH}$ and the velocity dispersion $\\sigma$ of the host galaxy (Ferrarese \\& Merritt 2000; Gebhardt et al. 2000; Merritt \\& Ferrarese 2001b) offers a new tool for evaluation of BH masses in AGNs if a reliable measurement of host galaxy luminosity or velocity dispersion is done.  While for AGN with strong emission lines (as QSO and Seyfert galaxies) the standard methods (e.g. reverberation mapping) under virial assumptions of the emitting regions can be  used to derive M$_{BH}$ ( see e.g. Wandel et al. 1999 and Kaspi et al. 2000), the above relations may be  the only way to estimate M$_{BH}$ for active galaxies that lack of emission lines (as BL Lac objects) or that are too far (most of nearby radiogalaxies) to resolve the region of influence of the BH.  In addition the knowledge of the properties of the galaxies hosting active nuclei may also yield fundamental insight for probing the unification models of AGN. The main ingredient of the unification scenario of radio loud AGN (e.g. Urry and Padovani 1995) is that the observed properties of an object may depend on orientation effects (because of an isotropic obscuration or emission from the nucleus). Two objects with identical intrinsic properties may therefore exhibit apparent different nuclear activity because they are seen at different angles.  A simple test for this hypothesis is  to compare the properties of the host galaxies (that do not depend on orientation) of various classes of AGN. Radio loud AGNs constitute only a small fraction of active galaxies but they seem to form a  homogeneous class of sources whose phenomenology could be explained within the same scenario.  In this paper I'll first review the properties of the galaxies hosting low redshift (z$<$ 0.5) radio loud active nuclei. This includes radiogalaxies, BL Lac objects and radio loud quasars (RLQs). Then using the relationships between BH mass and host properties I  derive and compare the distribution of BH masses in the various classes of active galaxies. In the second part of this work I'll sketch the view of quasar hosts beyond z $\\sim$ 1. Finally a summary of the main conclusions of this work are given together with a perspective of the future studies on high redshift sources. ",
        "conclusions": "I have shown that at low redshift (z $<$ 0.5)the galaxies hosting radio sources exhibit very similar properties. In all classes considered (RG, BL Lacs and RLQs) the hosts are luminous galaxies dominated by the spheroidal component. These objects seem to follow the behaviour of massive spheroids that are well formed at z $>$ 2 and then undergo passive evolution.  Assuming that the relationship between galaxy mass (luminosity or velocity dispersion) and central black hole mass, found for nearby early type galaxies, holds also for these more distant active galaxies, the similarity of host properties translates into a similarity of M$_{BH}$. The differences in the observed nuclear properties must therefore be  either in their viewing angle or/and  in the level of accretion. Also a different BH spin could play a relevant role.  Further studies of the host galaxies and their nuclei require to extend the analysis to higher redshift. In particular it is important to explore, with a statistical data set, QSO hosts up to z $\\sim$ 3 in order to assess the host evolution around the peak of QSO activity. To pursue this goal it is imperative to gather observations with high spatial resolution and great efficiency in order to detect the faint  nebulosity surrounding the bright nuclei. These requirements appear well matched by next generation  instrumentation, in the near IR, that make use of adaptive optics at large ground based telescopes or with the future James Webb Space Telescope  in space."
    },
    "0212/astro-ph0212487_arXiv.txt": {
        "abstract": "{ We have identified a new early T dwarf only 3.6\\,pc from the Sun, as a common proper motion companion (separation 1459\\,AU) to the K5V star \\thestar{} (HD\\,209100). As such, \\thetdwarf{} is one of the highest proper motion sources outside the solar system ($\\sim$\\,4.7 arcsec/yr), part of one of the twenty nearest stellar systems, and the nearest brown dwarf to the Sun. Optical photometry obtained from the SuperCOSMOS Sky Survey was combined with approximate infrared photometry from the 2MASS Quicklook survey data release, yielding colours for the source typical of early T dwarfs. Follow-up infrared spectroscopy using the ESO NTT and SOFI confirmed its spectral type to be T2.5$\\pm$0.5. With $K_s$=\\magnit{11}{2}, \\thetdwarf{} is 1.7 magnitudes brighter than any previously known T dwarf and 4 magnitudes brighter than the typical object in its class, making it highly amenable to detailed study. Also, as a companion to a bright nearby star, it has a precisely known distance (3.626\\,pc) and relatively well-known age (0.8--2\\,Gyr), allowing us to estimate its luminosity as log\\,L/\\Lsolar=$-4.67$, its effective temperature as 1260\\,K, and its mass as $\\sim$\\,40--60\\Mjup{}. \\thetdwarf{} represents an important addition to the census of the Solar neighbourhood and, equally importantly, a new benchmark object in our understanding of substellar objects. ",
        "introduction": "Our knowledge of the solar neighbourhood remains remarkably sketchy. Recent discoveries of very nearby ($<$10\\,pc) early and mid M dwarfs (Scholz, Meusinger, \\& Jahrei{\\ss} \\cite{scholz01}; Reid \\& Cruz \\cite{reid02a}; Reid, Kilkenny, \\& Cruz \\cite{reid02b}; Reyl\\'{e} \\etal{} \\cite{reyle02}; L\\'{e}pine, Rich, \\& Shara \\cite{lepine02a}) and late M dwarfs (Delfosse \\etal{} \\cite{delfosse01}; McCaughrean, Scholz, \\& Lodieu \\cite{mccaughrean02}; L\\'{e}pine \\etal{} \\cite{lepine02b}) appear to vindicate predictions that up to 30\\% of all stars within 10\\,pc remain unknown (Henry \\etal{} \\cite{henry97}).  The deficiency is even larger at substellar masses. Observational evidence (Reid \\etal{} \\cite{reid99}) and theory (Chabrier \\cite{chabrier02}) suggest a local space density of $\\sim$0.1 per cubic parsec for brown dwarfs, \\idest{} twice that for main sequence stars. However, brown dwarfs are much cooler and fainter and thus harder to observe, and as a consequence, only one M9 brown dwarf (Tinney \\cite{tinney96}, \\cite{tinney98}), a handful of L dwarfs, and 10 T dwarfs (see Burgasser \\etal{} \\cite{burgasser03} and references therein) are known at less than 10\\,pc from the Sun, compared to $>$300 stars within the same volume.  These sources are important however, as the detailed observation of very nearby stars and brown dwarfs is a key starting point for investigations of the star formation process, the stellar and substellar luminosity function, and the initial mass function. The recently discovered and categorised T dwarfs, with temperatures of $\\sim$\\,1500\\,K or less (Burgasser \\etal{} \\cite{burgasser02}) are particularly interesting, as they span a range of masses from just substellar down close to those of giant planets. Some 32 spectrally confirmed T dwarfs have been found to date, most of them in wide-field infrared (2MASS; Burgasser \\etal{} \\cite{burgasser02}) and optical (SDSS; Geballe \\etal{} \\cite{geballe02}) surveys. However, as substellar brown dwarfs continuously cool and decrease in luminosity throughout their lifetime, it is relatively difficult to determine their distances, ages, and masses. There are two important exceptions, Gl229\\,B and Gl570\\,D, both of which are companions to nearby stars for which distances and ages are relatively well known, and these key template objects help underpin studies of brown dwarf evolution, chemistry, and weather. It is crucial that we expand the sample of T dwarfs with well-understood physical properties, and the nearer they are to the Sun, the better, as they can be studied in more detail. In this paper, we describe the discovery of a T dwarf as as a companion to one of the nearest stars in the sky, \\thestar. ",
        "conclusions": "Through a survey for high proper motion objects, we have identified \\thetdwarf, a T2.5 dwarf, the brightest object in the T spectral class, and the nearest bona fide brown dwarf to the Sun. In combination with its accurately known distance, relatively well-known age, and large separation from its primary star, these characteristics make \\thetdwarf{} a new benchmark in the study of substellar objects, amenable to a wide range of detailed atmospheric and chemical observations. Urgently needed, however, are more accurate near-infrared photometry, including measurements out to the $L$ band, and spectral model fitting to derive an explicit bolometric magnitude for \\thetdwarf, and thus help calibrate the presently poorly-understand L-T dwarf boundary."
    },
    "0212/astro-ph0212074_arXiv.txt": {
        "abstract": "We model the Class I source L1551 IRS 5, adopting a flattened infalling envelope surrounding a binary disk system and a circumbinary disk. With our composite model, we  calculate self-consistently the spectral energy distribution of each component of the  L1551 IRS 5 system, using additional constraints from recent observations by ISO, the water ice feature from  observations with  SpeX, the SCUBA extended spatial brightness distribution at sub-mm wavelengths, and the VLA spatial intensity distributions at 7 mm of the binary disks. We analyze the sensitivity of our results to the various parameters involved.  Our results show that a flattened envelope collapse model is required to explain simultaneously the large scale fluxes  and the water ice and silicate features. On the other hand, we find that the circumstellar disks are optically thick in the millimeter range and are inclined so that their outer parts hide the emission along the line of sight  from their inner parts.  We also find that these disks have lower mass accretion rates than  the infall rate of the envelope. ",
        "introduction": "\\label{introdu}  L1551 IRS 5 in Taurus (distance 140 pc; Kenyon, Dobrzycka \\& Hartmann 1994) is one of the most extensively studied young stellar objects (YSOs). It is the prototypical embedded YSO (Strom, Strom, \\& Vrba 1976) with a massive molecular outflow (Snell, Loren \\& Plambeck et al. 1980). Adams, Lada, \\& Shu (1987) identified it as a protostellar (Class I)  source with an estimated infall rate for its envelope of $\\sim 10^{-5} \\msunyr$. The near-infrared spectrum exhibits first-overtone absorption in CO, suggesting that IRS 5 is an FU Ori object (Carr, Harvey, \\& Lester 1987), consistent with an analysis of the optical spectrum observed in scattered light (Mundt et al. 1985 and Stocke {\\etal} 1988).  Given the wide variety of observational constraints available, L1551 IRS 5 is clearly an ideal system for investigating the circumstellar environment that surrounds protostars and to understand its implications for the general process of star formation.  The spectral energy distribution (SED) of dust thermal emission in L1551 IRS 5 and spatial intensity profile at several wavelengths have been modeled previously by several authors (Adams, Lada, \\& Shu 1987; Butner et al. 1991; Kenyon, Calvet \\& Hartmann 1993, hereafter KCH93; Men'shchikov \\& Henning 1997;  White et al. 2000). Many of the models developed so far predict SEDs with silicate absorption features much deeper than that revealed by recent observations with the Infrared Space Observatory (ISO) and adopt outer envelopes smaller than indicated by sub-mm SCUBA observations (Chandler \\& Richer 2000). Some models assume simple power-law density distributions for the circumstellar envelope without requiring physical self-consistency.  In addition, in most of the models, L1551 IRS 5 is considered as a single-star system, while high-resolution interferometric observations (Rodr\\'\\i guez et al. 1998) indicate that IRS 5 is a binary.  In this work, we present a physically self-consistent model to investigate the circumstellar material in L1551 IRS 5 on different scales. We assume flattened envelope models (Hartmann et al. 1994;  Hartmann, Calvet \\& Boss 1996, HCB96) for which we solve the transfer of radiation following the methods described by KCH93 and HCB96. We further adopt viscous accretion disks that take into account the irradiation by the envelope such as those described by D'Alessio, Calvet \\& Hartmann (1997, DCH97).  With this composite model, we self-consistently calculate the overall SED and spectral features, the sub-mm spatial intensity profile, the SEDs of the inner disks, and compare with observations to constrain the envelope and disk parameters.  This paper is organized as follows. In \\S 2 we present a description of the L1551 IRS 5 system, with a summary of the observations used in our modeling efforts.  In \\S 3 we report spectra at 2-5 microns which constrain disk properties and ice absorption. In \\S 4 we discuss the main assumptions of our modeling.  In \\S 5 we present our results and in \\S 6 we discuss them. Finally, in \\S 7 we present a summary and conclusions. ",
        "conclusions": "Our main results can be summarizes as follows:  \\begin{enumerate}  \\item The wealth of observational data available for L1551 IRS 5, in particular the observed SED  by ISO and the water ice feature  at 3 microns by SpeX has enabled us to develop detailed models for the circumstellar envelope,  constraining its luminosity, geometry, infall accretion  rate, and  opacity.  \\item We find that flattened collapse models with an inclination along the line of sight $\\sim 50^{\\circ}$, a total luminosity of $\\sim 25~L_{\\odot}$, mass $\\sim 4~ M_{\\odot}$, infall accretion rate  $\\sim$ 7$\\times 10^{-5}$ $M_{\\odot}~ yr^{-1}$ and a normal opacity laws, modified by changing the water ice abundance, enable us to fit simultaneously the ISO, SpeX  observations and  single dish submillimeter and millimeter fluxes that are mostly dominated  by the emission of the large scale envelope.  \\item We  find that the a mixture as the proposed by Pollack et al. 1994 with no graphite but organic grains was not able to fit simultaneously the near-infrared and millimeter range of the spectrum. While we find reasonable fits to the envelope SED with the modified Draine \\& Lee opacities, we do not  conclude that graphite must necessarily be present, merely that the overall opacity curve must resemble the results for this mixture.   \\item  At smaller  scales than 2000 AU, the millimeter spatial intensity profiles in the SCUBA images suggest  the presence of the circumbinary disk. We model this structure with a grain size distribution similar to that in the extended envelope and that in the circumstellar disks, as the two extreme possibilities for the state of the dust in this disk. We find that the mass of the circumbinary disk goes from 0.02 to 0.4 $M_{\\odot}$ for $a_{max}=400$ $\\mu$m and $a_{max}=0.3$ $\\mu$m, respectively.    \\item The properties of the circumstellar disks are inferred from high angular millimeter images and fluxes, optical spectra and near-IR image. We find that these disks are optically thick in the mm wavelength range; and  are inclined  an angle ($\\sim 60^{\\circ}$), so that  their outer parts hide  the emission from their inner parts. We also find that these disks have larger mass accretion rates ($\\sim  2 \\times 10^{-6} \\ \\msunyr$) than the typical Classical T Tauri disks, usually optically thin at radio frequencies.   \\item The millimeter emission of the circumstellar disks emerges from a depth close to the midplane. As a consequence of this, the disk emission is characterized by a temperature much higher than the disk photospheric temperature. This explains why the disks have a brightness temperature higher than expected for the photosphere of a viscous disk, even with a high mass accretion rate   and illustrates the importance of constructing detailed physical vertical structure models reflecting the actual temperature gradients, and also the power of interferometric imaging in the mm range to probe the disk interior.  \\item The circumstellar disk models with parameters that reproduce the observational constraints, predict very well the observed strengths of the first two CO first overtone bands detected in our SpeX spectrum, providing an additional confirmation of the disk structural parameters.   \\item In the case of L1551 IRS 5, we have for the first time the confirmation of the suggestion, that material is piling up onto outer disks without immediately and continuously accreting inward since the disk accretion rate directly estimated from the mm observations is much lower than the infall rate, probably resulting in the accumulation of material in a circumbinary disk, possibly with occasional cascades of accretion into the binary system to produce FU Orionis eruptions.   \\end{enumerate}"
    },
    "0212/gr-qc0212033_arXiv.txt": {
        "abstract": "We consider the application of the consistent lattice quantum gravity approach we introduced recently to the situation of a Friedmann cosmology and also to  Bianchi  cosmological models.  This allows us to work out in detail the computations involved in  the determination of the Lagrange multipliers that impose consistency, and the implications of this determination.  It also allows us to study the removal of the Big Bang singularity.  Different discretizations can be achieved depending on the version of the classical theory chosen as a starting point and their relationships studied. We analyze in some detail how the continuum limit arises in various models. In particular we notice how remnants of the symmetries of the continuum theory are embodied in constants of the motion of the consistent discrete theory. The unconstrained nature of the discrete theory allows the consistent introduction of a relational time in quantum cosmology, free from the usual conceptual problems. ",
        "introduction": "The use of lattice regularizations in quantum gravity has always been problematic due to the fact that the introduction of the lattice breaks diffeomorphism invariance (see Loll \\cite{loll} for a review). The lattice theory therefore has a very different symmetry group than the continuum theory it approximates. We have recently introduced \\cite{us} a procedure to treat in a consistent way the resulting discrete theories, including their symmetries, both at a classical and quantum mechanical level. As is common in discrete theories, the resulting theory fixes Lagrange multipliers that in the continuum theory are free \\cite{TDLee}. This is in principle puzzling and it is worthwhile examining in detail in situations where one can easily control all aspects. In this paper we study these issues in the context of  cosmological models: the Friedmann and the Bianchi models. These models have several appealing properties: they have enough degrees of freedom as to have a set of non-trivial observables and make the deparameterization and physical description of the evolution possible, which is a feature one would like to analyze in the full theory. The Bianchi II model is classically soluble through a canonical transformation that maps it to a free particle. One can discretize the model before or after the canonical transformation. The discrete theories are not the same and there is the question as of their similarities and differences. Moreover, generically, other canonical transformations (like the one that introduces the Ashtekar variables) in the continuum will also lead to different theories upon discretization. We will see that the discrete theories avoid the singularity and have symmetries that translate themselves in constants of the motion that are remnants of the observables of the continuum theory. We will see that the discrete theory allows naturally to introduce a relational notion of time free of the usual problems of such approach.    We have recently introduced \\cite{us} a general technique for treating the theories that arise when one discretizes a continuum field theory. We have shown that the technique works for Yang--Mills and BF theories and implemented it for the gravitational case. Here we summarize the basic features of the approach. Since lattice field theories have a finite number of degrees of freedom, and since in this paper we are looking at cosmologies only, it suffices to summarize for illustration purposes the technique when applied to a constrained mechanical system.  Consider a constrained mechanical system. We assume we replace in the action all time derivatives with discrete first order finite differences. The time integral in the action is replaced by a sum $S=\\sum_{n=1}^N L(q_n,q_{n+1})$ with, \\begin{equation} L(n,n+1) = p_n (q_{n+1}-q_n) -\\epsilon H(q_n,p_n)-\\lambda_{nB} \\phi^B(q_n,p_n) \\end{equation} where we assume we have a Hamiltonian $H$ and $M$ constraints $B=1\\ldots M$ and $\\epsilon$ is the time interval. Our formulation works for a system with an arbitrary (finite) number of phase-space degrees of freedom. To simplify the notation we are omitting the indices labeling the degrees of freedom and writing formulae as if the system only had one configuration degree of freedom.  The centerpiece of our technique is to realize that in a theories where time (and space in the case of field theories) is discretized, it makes no sense to work with a Hamiltonian since the latter is the generator of infinitesimal time evolution and one cannot have infinitesimal evolution if time is discrete. The time evolution should be described via a canonical transformation that implements the discrete time evolution between instants $n$ and $n+1$. A type 1 canonical transformation that accomplishes this task has as generating function minus the Lagrangian, viewed as a function of $q_n$ and $q_{n+1}$. From it, one obtains the momenta at instant $n+1$ \\begin{eqnarray} P^q_{n+1} &=& {\\partial L(n,n+1) \\over \\partial q_{n+1}} =p_n\\label{37}\\\\ P^p_{n+1}&=& {\\partial L(n,n+1) \\over \\partial p_{n+1}}= 0\\label{38}\\\\ P^{\\lambda_B}_{n+1}&=& {\\partial L(n,n+1) \\over \\partial \\lambda_{(n+1)B}}= 0, \\end{eqnarray} and similarly, the momenta at instant $n$ are given by, \\begin{eqnarray} P^q_n &=& -{\\partial L(n,n+1) \\over \\partial q_{n}} =p_n+\\epsilon {\\partial H(q_n,p_n) \\over \\partial q_n} + \\lambda_{nB} {\\partial \\phi^B(q_n,p_n) \\over \\partial q_n}\\label{40}\\\\ P^p_n &=& -{\\partial L(n,n+1) \\over \\partial p_{n}} = -(q_{n+1}-q_n) + \\epsilon {\\partial H(q_n,p_n) \\over \\partial p_n} + \\lambda_{nB} {\\partial \\phi^B(q_n,p_n) \\over \\partial p_n}\\label{41}\\\\ P^{\\lambda_B}_n &=& \\phi^B(q_n,p_n).\\label{42} \\end{eqnarray} We can now combine these equations to give, \\begin{eqnarray} p_n-p_{n-1} &=& -\\epsilon {\\partial H(q_n,p_n) \\over \\partial q_n} - \\lambda_{nB} {\\partial \\phi^B(q_n,p_n) \\over \\partial q_n}\\label{43}\\\\ q_{n+1}-q_n &=& \\epsilon {\\partial H(q_n,p_n) \\over \\partial p_n} + \\lambda_{nB} {\\partial \\phi^B(q_n,p_n) \\over \\partial p_n}\\label{44}\\\\ \\phi^B(q_n,p_n) &=&0.\\label{45} \\end{eqnarray} Superficially, these equations appear entirely equivalent to the continuum ones. However, they hide the fact that in order for the constraints to be preserved, the Lagrange multipliers get fixed. Another way to see it, is that $ P^q_{n+1} =p_n$ and therefore it is immediate that the Poisson bracket of the constraints evaluated at $n$ and at $n+1$ is non-vanishing. It is illuminating to rewrite these equations in terms of the canonically conjugate pair, \\begin{eqnarray} P^q_{n+1}-P^q_{n} &=& -\\epsilon {\\partial H(q_n,P^q_{n+1}) \\over \\partial q_n} - \\lambda_{nB} {\\partial \\phi^B(q_n,P^q_{n+1}) \\over \\partial q_n}\\label{11}\\\\ q_{n+1}-q_n &=& \\epsilon {\\partial H(q_n,P^q_{n+1}) \\over \\partial P^q_{n+1}} + \\lambda_{nB} {\\partial \\phi^B(q_n,P^q_{n+1}) \\over \\partial P^q_{n+1}}\\label{12}\\\\ \\phi^B(q_n,P^q_{n+1}) &=&0.\\label{13} \\end{eqnarray}  We therefore will eliminate the ``constraints'' (\\ref{13}) by solving them for the Lagrange multipliers. One can proceed in various ways, either determining the Lagrange multipliers at time $n$ as functions of the variables at either $n,n-1$ or $n+1$.  We will choose to determine the Lagrange multipliers at instant $n$ as  functions of $P^q_n$ and $q_n$ in this paper. For different problems it can be more convenient to make one choice over the other. In all cases solving for the Lagrange multipliers implies solving algebraic equations, but depending on the choices made the resulting equations can be quite non-linear and guaranteeing that they will have real solutions can be problematic. For the examples of this paper we have found that the choice we make is the most convenient one.  We solve (\\ref{11}) for $P^q_{n+1}$ and substitute it in the ``constraint'' (\\ref{13}) and obtain \\begin{equation} \\phi^B(q_{n},P^q_{n},\\lambda_{nB})=0, \\label{system} \\end{equation} and this constitutes a system of equations. Generically, these will determine \\begin{equation} \\lambda_{nB} =\\lambda_{nB}(q_{n},P^q_{n}, v^\\alpha) \\label{lambdadet} \\end{equation} where the $v^\\alpha$ are a set of free parameters that may arise if the system of equations is undetermined. The eventual presence of these parameters will signify that the resulting theory still has a genuine gauge freedom represented by freely specifiable Lagrange multipliers.  The final set of evolution equations for the system is therefore given by (\\ref{11},\\ref{12}) where the Lagrange multipliers are substituted using (\\ref{lambdadet}).    This construction raises several questions about how to implement it in concrete gravitational examples. We would like to probe these questions in the context of cosmological models. The main points we want to probe are the following:  {\\em (i) Solubility of the multiplier equations:} Solving the constraints by choosing the Lagrange multipliers produces a theory that is constraint free. This is analogous to what happens when one gauge fixes a theory. It is well known that gauge fixing is not a cure for the problems of general relativity since gauge fixings usually become problematic. In the same sense, it could happen that the algebraic equations that determine the Lagrange multipliers (the lapse and the shift in the case of general relativity) in our approach develop problems in their solutions (for instance, negative lapses, or complex solutions). We will show that in these simple examples there do not appear to be difficulties even in situations where a traditional gauge fixing in the continuum theory is not possible.  {\\em (ii) Performing meaningful comparisons:} When comparing a discrete theory with a continuum theory, one needs to choose the quantities that are to be compared. In particular, the continuum theory has ``observables'' (``perennials''), that is, quantities that have vanishing Poisson brackets with the constraints.  The discrete theory is constraint-free. We will see that one has exact conserved quantities in the discrete theory that are consequences of remnant symmetries encoded in the initial data that the discrete theory inherits from the continuum theory. The conserved quantities will turn out to be closely related to the observables of the continuum theory.   {\\em (iii) The continuum limit}: in continuum constrained theories with first class constraints, the Lagrange multipliers are free functions. Yet in our discrete construction, the Lagrange multipliers are determined by the initial conditions. If one wishes to take a naive continuum limit, the discrete equations that determine the Lagrange multipliers must become ill-defined. To extract meaningful information about the continuum theories, one needs to proceed differently. We have already mentioned that the discrete theories are constraint-free, since one solves the constraints for the Lagrange multipliers. This in particular means that they have more degrees of freedom than the continuum theory they are trying to approximate. The continuum limit might be achieved by a careful fine tuning of the extra degrees of freedom, or might occur in certain asymptotic regions of the dynamics as an attractor for a large set of initial values of the extra degrees of freedom. We will exhibit examples of both kinds of behaviors.    {\\em (iv) Singularities:} Discrete theories in principle have the possibility of evolving through a Big Bang singularity, since generically the singularity will not lie on a point on the lattice. However, we will see that if one uses canonical variables such that the singularity is at a boundary of the range of a variable, then the discrete theories do develop singularities, although they can be avoided in certain cases.  {\\em (v) Problem of time:} We will discuss how to obtain evolution in the discrete by using a relational approach in terms of the observables of these theories. This suggests that the problem of time can be solved in these theories.   {\\em (vi) Discretization ambiguities}: An important element is to note that the Lagrange multipliers get determined by this construction only if the constraint is both a function of $q$ and $p$. If the constraint is only a function of $q$ or of $p$ then the constraints are automatically preserved in evolution without fixing the Lagrange multipliers. This raises a conceptual question.  For certain theories in the continuum, one can make a canonical transformation to a new set of variables such that the constraints depend only on $q$ or on $p$. The resulting discrete theories will therefore be very different in nature, but will have the same continuum limit. From the point of view of using discrete theories to quantize gravity, we believe this ambiguity should receive the same treatment that quantization ambiguities (choice of canonical variables, factor orderings, etc.)  get: they are decided experimentally. Generically there will be different discrete theories upon which to base the quantization and some will be better at approximating nature than others in given regimes. Many of them may allow to recover the continuum limit, however they may have different discrete properties when one is far from the semiclassical regime.  The organization of this paper is as follows: in the next section we discuss several Friedmann models that exhibit various behaviors upon discretization. In section III we discuss the quantization of the models and the various possibilities concerning how one approaches the ``problem of time''. In section IV we discuss other models, including the Bianchi ones and show that they exhibit similar kinds of behaviors as the ones we encountered in the Friedmann models. We end with a discussion of the importance of the results. ",
        "conclusions": "We have studied the recently introduced ``consistent discretization'' approach to general relativity in the case of cosmological models. This helps to clarify, through examples, several questions that can be raised concerning that approach. First and foremost, the unusual property of having the Lagrange multipliers determined operates without difficulty in the models considered.  We never encounter that the equations imply that the lapse is complex or that it is multi-valued during evolution. Moreover, the issue of how the discrete theory achieves a continuum limit is illuminated. We see that the possibility of having a continuum limit depends on the type of model and of discretization chosen. Some models exhibit continuum behavior spontaneously as an attractor in a region of their evolution. In other models, careful tuning of initial data is needed to approximate the continuum theory. The Big Bang singularity can be avoided in a generic fashion in some models, whereas in others it appears as inevitable, and this can be forecasted from the way the singularity appears in the continuum theory. We also saw explicitly how different formulations of the continuum theory can, upon discretization, exhibit radically different behaviors at the level of the discrete theory. Finally, the quantization of the discrete theory, in spit of the fact that the latter has no constraints, is non-trivial. One discovers that the discrete theories have remnant symmetries that imply the existence of constants of the motion that are remnants of the observables (``perennials'') of the continuum theory. This opens several possibilities for the quantization, relating to how one approaches the ``problem of time''. One can attempt a straightforward Schr\\\"odinger quantization of the theory, but the evolution one obtains is in terms of a ``time'' that has no physical meaning. One can attempt to ``gauge fix'' the remnant symmetry and one obtains a theory with a physically well motivated time that can be quantized a la Schr\\\"odinger, but it has the drawback that one has artificially decided to treat a variable as classical, even in regimes in which it should not be. Finally, one can attempt a relational introduction of the notion of time that is free of the many problems that plagued such attempts in the continuum theory and therefore provides a natural solution to the problem of time in quantum cosmology.  It is clear that one can only take the examples presented in this paper, given their simplicity, as first attempts to understand this approach to discrete general relativity. Further work will be needed to gain confidence that the approach can be used in more realistic settings. But already the exploration of this simple models has been a useful laboratory to exhibit the kind of behaviors that one may encounter in the discretized theories."
    },
    "0212/astro-ph0212242_arXiv.txt": {
        "abstract": "The build--up of stellar mass in galaxies is the consequence of their past star formation and merging histories.  Here we report measurements of rest--frame optical light and calculations of stellar mass at high redshift based on an infrared--selected sample of galaxies from the Hubble Deep Field North.   The bright envelope of rest--frame $B$--band galaxy luminosities is similar from $0 < z < 3$, and the co--moving luminosity density is constant to within a factor of 3 over that redshift range.  However, galaxies at higher redshifts are bluer, and stellar population modeling indicates that they had significantly lower mass--to--light ratios than those of present--day $L^\\ast$ galaxies.  This leads to a global stellar mass density, $\\Omega_\\ast(z)$, which rises with time from $z = 3$ to the present.  This measurement essentially traces the integral of the cosmic star formation history that has been the subject of previous investigations.  50--75\\% of the present--day stellar mass density had formed by $z \\sim 1$, but at $z \\approx 2.7$ we find only 3--14\\% of today's stars were present.  This increase in $\\Omega_\\ast$ with time is broadly consistent with observations of the evolving global star formation rate once dust extinction is taken into account, but is steeper at $1 < z < 3$ than predicted by some recent semi--analytic models of galaxy formation.  The observations appear to be inconsistent with scenarios in which the bulk of stars in present--day galactic spheroids formed at $z \\gg 2$. ",
        "introduction": "}  One goal of observational cosmology is to understand the history of mass assembly in galaxies.  In current models of structure formation, dark matter halos build up in a hierarchical process controlled by the nature of the dark matter, the power spectrum of density fluctuations, and the parameters of the cosmological model.  The assembly of the stellar content of galaxies is governed by more complex physics, including gaseous dissipation, the mechanics of star formation itself, and the feedback of stellar energetic output on the baryonic material of the galaxies.  While it is clear that the mean stellar mass density of the universe, $\\Omega_\\ast$, should start at zero and grow as the universe ages, the exact form of this evolution is not trivial, and could reveal much about the interaction of large--scale cosmological physics and small--scale star formation physics.  For example, there may have been multiple epochs of star formation in the universe, associated with changes in the ionizing radiation background.  The initial mass function (IMF) may have been different for early generations of stars or in different environments (e.g., in starbursts versus quiescent disks). This would change the relation between luminosity, metal production and total star formation rates.  Early star formation with a top--heavy IMF might leave little in the way of remaining stellar mass, while contributing significantly to the ionization of the intergalactic medium and dispersing metals into it.   The total, integrated mass in stars is tightly coupled to the history of star formation as traced by the infrared background, and to the cold gas content of the universe, $\\Omega_{\\mathrm{HI}}$.  Reducing the uncertainties on all of these measurements will provide strong constraints on models for galaxy formation, while discrepancies could reveal new physics.  Deep imaging and spectroscopic surveys now routinely find and study galaxies throughout most of cosmic history, back to redshifts $z = 3$ and earlier.  A variety of studies have considered the evolution of the global star formation rate (i.e., averaged over all galaxies), using different observational tracers of star formation in distant galaxies \\citep{lil96, mad96, mad98, ste99, bar00, tho01}, or constraints from neutral gas content and/or the extragalactic background light \\citep{pei95, pei99, cal99}.  It is generally accepted that star formation in galaxies was more active at high redshift than today. Several investigations have found an approximately constant rate of global star formation from $z \\approx 4$ to 1, with a steeper decline since that epoch to the present--day, although other studies have questioned these results \\citep{cow99, wil02, lan02}.  The true rates of star formation (either globally or for individual galaxies) are subject to a variety of uncertainties and biases, particularly regarding dust extinction and its effect on galaxy selection and on observational indicators of star formation.  A complementary approach is to try to measure the stellar {\\it masses} ($\\mathM$) of galaxies, rather than the {\\it rates} at which they are forming stars.  This is, at least indirectly, a common motivation for near--infrared galaxy surveys, which observe starlight that provides a reasonable tracer of the total stellar mass.  The stellar mass of a galaxy is dominated by older, lower luminosity, redder stars. At longer wavelengths, these stars provide a greater fractional contribution to the integrated light from the galaxy, relative to that from short--lived, massive stars.  For a stellar population formed over an extended period of time, the luminosity--weighted mean age increases at redder wavelengths, and the range of possible mass--to--light ratios ($\\mathM/L$) is reduced.  Moreover, the impact of dust extinction is greatly reduced at longer wavelengths.  Measurements of the $K$--band luminosity function of nearby galaxies have been aimed, in part, at assessing the distribution of stellar mass among galaxies at the present day. For example, in a recent study, \\citet{col01} combined data from the 2dF Galaxy Redshift Survey (2dFGRS) and the Two Micron All--Sky Survey (2MASS) to produce a large and well--controlled sample for determining the $K$--band luminosity function.  They employed stellar population synthesis models to relate light to mass, using the colors of galaxies to constrain $\\mathM/L$, and hence to derive the stellar mass function of galaxies in the nearby universe.   \\citet{kau02} have recently used $z$--band (0.9$\\mu$m) photometry and spectroscopic stellar population indices in order to evaluate the stellar mass content of a large sample of galaxies from the Sloan Digital Sky Survey (SDSS).  Similar methods have been applied to galaxies at cosmological distances.  Deep near--infrared surveys are generally intended to measure the rest--frame optical light from high redshift galaxies, in order to minimize the effects of active star formation and dust on galaxy selection and to more nearly trace their stellar masses. \\citet{bri00} have used $K$--band photometry of galaxies from redshift surveys of {\\it Hubble Space Telescope (HST)} imaging fields to derive stellar masses for galaxies out to $z \\approx 1$, studying the evolution of the global stellar mass density as a function of galaxy morphology.  \\citet{coh02} has also combined extensive redshift survey data with infrared photometry and evolutionary models to study the evolving stellar mass density out to $z \\approx 1$.  \\citet{dro01} have used an infrared imaging and photometric redshift survey, which covers a larger solid angle to shallower depths, to study the distribution function of stellar mass at the bright end of the galaxy population over a similar redshift range. At higher redshift, \\citet{saw98}, \\citet{pap01} and \\citet{sha01} have used deep infrared data to study the stellar mass content of Lyman Break Galaxies (LBGs;  cf., Steidel \\etal\\ 1996), actively star--forming galaxies at $z \\sim 3$ that are selected on the basis of their UV luminosity and the broad band color signature of the Ly~$\\alpha$ and Lyman limit breaks.  Here, we will follow the basic approach of \\citet{col01}, \\citet{bri00} and \\citet[henceforth PDF]{pap01} to determine stellar masses for galaxies in the Hubble Deep Field North (HDF--N) out to $z = 3$, and to study the redshift evolution of the global stellar mass density. Throughout this paper, we assume a cosmology with $\\Omega_{\\rm tot}, \\Omega_M, \\Omega_\\Lambda = 1.0, 0.3, 0.7$, and $h \\equiv H_0 / (100~\\mathrm{km~s^{-1}~Mpc^{-3}}) = 0.7$. We will sometimes write $h_{70} \\equiv H_0 / (70~\\mathrm{km~s^{-1}~Mpc^{-3}})$. We use AB magnitudes ($AB \\equiv 31.4 - 2.5\\log\\langle f_\\nu / \\mathrm{nJy} \\rangle$), and denote the bandpasses used for {\\it HST} WFPC2 and NICMOS imaging as \\wfu, \\wfb, \\wfv, \\wfi, \\nicj, and \\nich. ",
        "conclusions": "}  We have used deep near--infrared imaging of the Hubble Deep Field North to select a sample of galaxies based on their optical rest--frame light out to $z = 3$, and have estimated their stellar mass content following the procedures outlined by \\citet{pap01}.  We find that the rest--frame $B$--band co--moving luminosity density is somewhat larger at high redshift than the present--day value, although constant to within a factor of 3 over the redshift range $0 < z < 3$.  However, galaxies at higher redshifts are much bluer than most luminous and massive galaxies locally, and the mean mass--to--light ratio evolves steeply with redshift.  The result is that the global stellar mass density, $\\Omega_\\ast$, rises rapidly with cosmic time, going from 3--14\\% of the present--day value at $z \\approx 2.7$, reaching 50--75\\% of its present value by $z \\approx 1$, and then (based on results from the literature) changing little to the present day.  In broad terms, the concordance between previous estimates of the cosmic star formation history and these new measurements of its time integral strongly suggests $1 < z < 2.5$ as a critical era, when galaxies were growing rapidly toward their final stellar masses. There are hints of a possible inconsistency between the global mass density $\\rho_\\ast$ at $z \\approx 3$ and evidence for high star formation rates at $3 < z < 6$.  One solution to this might be to tilt the IMF toward high--mass star formation at early times, as suggested by \\citet{fer02}.  \\citet{ren99} has recapped the evidence and arguments which suggest that spheroids contain $\\gtrsim 30$\\% of the stars in the local universe, and that those stellar populations are old and formed at high redshift. Recent estimates from the SDSS place $\\gtrsim 50\\%$ of the local stellar mass density in galaxies with red colors and/or large 4000\\AA\\ break amplitudes \\citep{hog02,kau02}.  Other estimates have set even larger values for the stellar mass fraction in galaxy spheroids (e.g., $\\sim 75\\%$ from \\citet{fuk98}).  This would appear to run counter to the evidence presented here, which suggests that $< 20$\\% of the present--day stellar mass was in place at $2 < z < 3$.  Our estimates of $\\Omega_\\ast(z)$ are not necessarily inconsistent with the ages of the red stellar populations which contribute so substantially to $\\Omega_\\ast$ today.  The presently favored world model dominated by a cosmological constant allows more time to elapse at lower redshifts for a given value of $H_0$.  In this case, the old ages of stars in local spheroids are more comfortably accommodated with lower formation redshifts, $z \\approx 2$, where the look--back time is $\\approx 10 h_{70}^{-1}$~Gyr.  We have found that most (at least half) of the present--day stellar mass was in place at $z \\approx 1$ (corresponding to a lookback time of $7.7h_{70}^{-1}$~Gyr), when the cosmic ``building boom'' seems to have been winding down.  It may be that the local amount of old, red starlight can be comfortably accommodated in such a universe.  A census of stellar populations in galaxies at $z \\approx 1$ may provide a stronger test of our HDF--N results at $z > 2$. \\citet{cim02a} have analyzed composite spectra of red galaxies at $z \\approx 1$, and have argued that a significant fraction of these galaxies are dominated by stellar populations with a minimum age of $\\sim 3$~Gyr.  Such stars therefore would have formed at $z \\gtrsim 2.3$ for the cosmology adopted here.  There is not yet a good estimate of the total, co--moving mass density in old stars at $z \\approx 1$, but if it were substantial this might prove to be discrepant with evidence for a rapid mass build--up from $z \\approx 3$ to 1.  If necessary, it might be possible to push our estimate of $\\Omega_\\ast$($z > 2$) upward toward 30\\% of the present--day value if each galaxy were assigned its maximum stellar mass (at, say, 68\\% confidence) allowed by our 2--component star formation models -- e.g., if indeed there were a very early epoch of rapid star formation preceding the redshift range studied here. However, once again we note that these models seem {\\it ad hoc} and somewhat unphysical.  The robustness of these results is limited by several factors. First, there is the small volume of the HDF--N, and the vulnerability of a single sight--line to the effects of large scale structure, which we have discussed in \\S\\ref{section:field2field}.  Our results at $z \\sim 1$ generally agree well with other estimates of the luminosity and stellar mass density that can be found in the literature.  At higher redshifts, there is some evidence for differences between the HDF--N and HDF--S sight lines.  It is not yet clear if the implied stellar mass densities at high redshift in these two fields disagree by more than our estimated uncertainties for the HDF--N alone.  Second, photometric data reaching to the $K$--band can sample only rest--frame optical wavelengths at $z \\sim 3$, where the estimates of $\\mathM/L$ depend strongly on assumptions about the stellar population metallicity and past star formation history.  Also, at the highest redshifts, even the NICMOS data do not reach far enough down the galaxy luminosity function to constrain the total luminosity or stellar mass densities as tightly as one would like, leaving significant uncertainty on these integrated quantities.  However, if we are missing most of the total stellar mass density at $2 < z < 3$ for this reason, then it must be found primarily in low--luminosity galaxies.  This would strongly contrast with the situation today and at intermediate redshifts, where the bright end of the luminosity function dominates the integrated stellar mass density.  Finally, we note that our choice of cosmology does affect the redshift dependence of $\\Omega_\\ast(z)$. In a flat, matter dominated universe with the same value of $H_0$, the luminosity densities would be 80\\% larger at $2.5 < z < 3.0$. However, the {\\it ratio} of densities between $z \\approx 1$ and $z \\approx 2.7$ would be change by only 15\\%, since the change in world model would affect all of the high redshift points.  Therefore, the general conclusion of a steep increase in $\\Omega_\\ast$ at $1 < z < 3$ would be essentially the same.  The launch of the {\\it Space Infrared Telescope Facility (SIRTF)} should help to tie down some of these loose ends.  In particular, the Great Observatories Origins Deep Survey (GOODS, \\citet{dic02}) incorporates a {\\it SIRTF} Legacy program which will collect the deepest observations with that facility at 3.6--24$\\mu$m, in part to address this very issue.  The GOODS program should help to improve the situation in several ways.  It will survey two independent sight--lines, covering a solid angle $60\\times$ larger than that considered here.  {\\it SIRTF} observations out to $8\\mu$m will sample rest--frame $2\\mu$m light from galaxies out to $z = 3$, and $1\\mu$m light to $z = 7$.  This should should provide much better constraints on stellar masses than do our current NICMOS and ground--based data, which only reach optical rest--frame light at high redshift.  {\\it SIRTF} is a comparatively small telescope (85~cm aperture), and its sensitivity to the faintest galaxies will be limited by collecting area and source confusion. GOODS is designed to robustly detect typical ($\\sim 10^{10} M_\\odot$) Lyman break galaxies at $z = 3$ (and to push to fainter luminosities and higher redshifts over a smaller solid angle in its ultradeep survey). But just as we found with our NICMOS and ground--based data here, estimates of the total, integrated mass density may be limited by poor constraints on the faint--end slope of the galaxy luminosity function at these large redshifts.  Ultimately, observations from the {\\it James Webb Space Telescope}, which will reach nJy sensitivities with observations out to at least $5\\mu$m, may be required to fully address this problem, and to push toward still higher redshifts and earlier epochs."
    },
    "0212/astro-ph0212132_arXiv.txt": {
        "abstract": "{We show that a {\\it helical} shear flow of a magnetized plasma may serve as an efficient amplifier of Alfv\\'en waves. We find that even when the flow is purely ejectional (i.e., when no rotation is present) Alfv\\'en waves are amplified through the transient, shear-induced, algebraic amplification process. Series of transient amplifications, taking place sequentially along the flow, may result in a {\\it cascade amplification} of these waves. However, when a flow is swirling or {\\it helical} (i.e., some rotation is imposed on the plasma motion), Alfv\\'en waves become subject to new, much more powerful shear instabilities. In this case, depending on the type of differential rotation, both usual and parametric instabilities may appear. We claim that these phenomena may lead to the generation of large amplitude Alfv\\'en waves and the mechanism may account for the appearance of such waves in the solar atmosphere, in accretion-ejection flows and in accretion columns. These processes may also serve as an important initial (linear and nonmodal) phase in the ultimate subcritical transition to MHD Alfv\\'enic turbulence in various kinds of astrophysical shear flows. ",
        "introduction": "Most of astrophysical objects involve different kinds of plasma flows. Recently it was fully realized that collective phenomena in flows with spatially inhomogeneous velocities ({\\it shear flows} hereafter referred as SF) are characterized by remarkable, so called ``nonmodal\" processes, related with non-self-adjointness of linear dynamics of perturbations. Namely, it was found that SF: exchange energy with sound waves (\\cite{bf92}); couple different collective modes with one another and lead to their mutual transformations (\\cite{crt96}); generate nonperiodic, vortical modes of collective behaviour (so called ``shear vortices\" \\cite{rcm98}) - which eventually may or may not acquire wave-like features; excite {\\it beat wave} phenomena both in neutral fluids (\\cite{rm97}) and plasmas (\\cite{prm98}).  These processes take place not only in neutral fluids (\\cite{bf92, rm97}), standard MHD (\\cite{crt96, rmb96, tvy01}) and electrostatic plasmas (\\cite{rcm98, vyt00, mms00}), but also in strongly magnetized plasmas with anisotropic pressure (\\cite{crt97}), electron-positron plasmas (\\cite{mmr97}) and dusty plasmas (\\cite{pkr00}). The possible role of these phenomena in astrophysical context was immediately realized and a number of astrophysical applications, including pulsar magnetospheric plasmas (\\cite{mmr97}), solar atmospheric phenomena (\\cite{prm98,rpm00}) and galactic gaseous disk dynamics (\\cite{rph99}) appeared within a span of years.  One of the main shortcomings of all these studies, stemming from the very design of nonmodal schemes, is that the description  is given in the wave number space (${\\bf k}-space$),  the knowledge about the  appearance of shear-induced phenomena in the real, physical space is lacking. Two more serious limitations of these investigations are related with the neglect of the back reaction of perturbations on the mean flow and with the omission of viscous dissipation effects. However recently a successful effort was made in the direction of spatial visualization of the shear-induced wave transformations (\\cite{bprr01}), became clear that shear-induced processes in SF are quite robust and easily recognizable, even in the presence of quite heavy dissipation.  All these processes were  studied predominantly for simple, plane-parallel flow geometries with linear velocity shear profiles. However, recently, a new method was developed (\\cite{mr99}), which allows a {\\it local} analysis of the dynamics of linearized perturbations in SF with arbitrarily complex geometry and kinematics ${\\bf U}(x,y,z)$. It was found that a slight deviation from the plane-parallel mode of motion brings into the game a variety of new, exotic and asymptotically persistent modes of collective behavior.  Still the systematic investigation of shear-induced processes in kinematically complex SF is a challenging task in a state of infancy. The original study has been limited to two-dimensional flow patterns of neutral fluids (\\cite{mr99}).  One begins to wonder whether those astrophysical systems, where kinematically complex modes of plasma motion are definitely present --- e.g., astrophysical jets (\\cite{f98}) or solar tornados (\\cite{pm98, vl99}) --- might sustain these collective phenomena and what kind of observational consequences could they lead to.  In this paper we consider interaction of a helical flow with Alfv\\'en waves generated within the flow. This problem is actual in plasma physics and plasma astrophysics from experimental/observational, theoretical and numerical points of view.  We are used to observe Alfv\\'en waves, being in interaction with plasma flows, in different astrophysical situations. For example, for the long time it is known  that the Sun radiates Alfv\\'en waves: outward propagating Alfv\\'en waves are routinely observed in the solar wind flow at $r>0.3\\,$AU (\\cite{h90}). The observations of quasiperiodic pulsed ionospheric flows (PIF) have shown that the PIF are driven and correlated with Alfv\\'enic fluctuations observed in the upstream solar wind (\\cite{ppmy02}). While the observations of the solar transition region (\\cite{p01}) imply that structuring of the transition region involves closed loops and coronal funnels showing unambiguous evidence for the presence of passing Alfv\\'en waves.  In the theoretical domain the problem of the interaction between Alfv\\'en waves and ambient flows is quite popular topic of studies in the wide range of applications including different kinds of laboratory, geophysical and astrophysical plasma flows. Recent studies of toroidal flows in axisymmetric tokamaks (\\cite{vbb00}), for instance, revealed that these flows generate low-frequency Alfv\\'en waves. The problem of Alfv\\'en waves sustained by plasma flows is especially popular in the context of solar physics. It is argued that chromosphere and transition region flows are primary energy sources for the fast solar wind; while Alfv\\'en waves, generated by post-reconnection processes, are continuously interacting with these flows (\\cite{rhwt01}). In accretion disks Alfv\\'en waves also seem to be actively interacting with the disk flow - they are found to be unstable both in high-$\\beta$ (\\cite{bh91}) and low-$\\beta$ (\\cite{tpc92}) disks. Most recently it was claimed that the Accretion-Ejection Instability (\\cite{tp99}) can extract accretion energy and angular momentum from magnetized disk, generate Alfv\\'en waves and ``feed\" the disk corona with them (\\cite{vt02}).  Currently sophisticated numerical codes are developed, which allow to study the propagation of Alfv\\'en waves along an open magnetic flux tube (\\cite{sks01}). These simulations are aimed to clarify mechanisms of the coronal heating and of the formation of solar plasma flow patterns: viz. solar spicules, macrospicules and solar tornadoes. More general numerical tools, like FINESSE and PHOENIX, aimed to study waves and instabilities in different kinds of flows, were very recently developed (\\cite{bbgvk02}). We assume that all these numerical tools could also be used in a number of astrophysical applications involving ``parent\" flows and ``inborn\" Alfv\\'en waves interacting with each other.  Most of above-cited studies considered simple (plane-parallel or locally plane-parallel) kinds of flows and the processes were treated by means of usual normal-mode analysis. The purpose of this paper (paper I) is to study `nonmodal' evolution of Alfv\\'en waves in swirling flows and to show that these flows may operate as efficient ``amplifiers\" of Alfv\\'en waves.  In the next section we develop the mathematical formalism and derive general equations governing the evolution of Alfv\\'en waves in helical flows. However, the third section is dedicated to the study of the simpler example of a parallel SF (`pure outflow\"). Still, even in this simple case, we find that the amplification of Alfv\\'en waves takes place. The amplification mechanism is linear, transient and can be described by appealingly simple mathematics. The amplification occurs within a relatively brief time interval and it appears as an abrupt, burst-like increase of the wave amplitude. It is tempting to argue that individual acts of wave amplifications, occurring sequentially, may ignite the ``chain reaction\" of {\\it nonmodal cascade} amplification of Alfv\\'en waves.  The fourth section of the paper is dedicated to fully helical SF. We show that when SF are helical (``outflow + rotation\") there appear new classes of shear instabilities capable of generating high-amplitude Alfv\\'en waves. The presence of differential rotation is crucial for these instabilities and in certain cases the instability is of parametric nature. These instabilities are powerful, often  have a `resonant' nature and are exclusively related with the kinematic complexity of `parent' shear flows.  The fifth section of the paper contains discussion of the obtained results and of their possible astrophysical applications. We argue that such wave amplification processes may have various astrophysical manifestations including large amplitude Alfv\\'en waves actually observed in the solar atmosphere (\\cite{bst95}). They could provide the necessary seeding for the development of MHD turbulence in hydromagnetic shear flows. These processes may be present and may lead to perceptible morphological variety and diverse observational appearances in various kinds of {\\it `accretion-ejection'} flows: innermost regions of accretion disks, disk-jet transition regions, accretion columns in X-ray pulsars and cataclysmic variables, inner regions of galactic gaseous disks, etc.  This paper will be followed by the second one (paper II), where we will consider the same flow structure but allow the perturbations to be fully compressible, bringing on-stage another two additional linear MHD wave modes - the slow and fast magnetosonic waves. ",
        "conclusions": "The results of the present study clearly show that both helical and purely ejectional flows of MHD plasmas could serve as efficient ``amplifiers\" of Alfv\\'en waves. In alliance with other shear-induced phenomena (such as shear-induced wave transformations and/or generation of shear vortices), these processes could lead to the cascade amplification of hydromagnetic waves leading to MHD turbulence in these  flows (\\cite{rpm00}). In more general terms  it could be stated that the presence of the velocity shear strongly modifies the nature of Alfv\\'en  waves such that they can extract energy from the mean flow.  Certainly, while talking about these processes, we should bear in mind that due to the limitation of the `nonmodal approach' we are not able to describe quantitatively the onset of these processes in the real physical space. Besides the linearized nature of our calculations prevents us from  taking into account of the back reaction of shear-modified Alfv\\'en waves on the mean flow. The last but not the least limitation is related with the neglect of the dissipation in the flow.  Can Alfv\\'enic perturbations still extract energy from the background flow when dissipation is present? For magnetoacoustic modes sustained by a simple, linearly sheared two-dimensional flow it was recently found (\\cite{bprr01}) that the presence of dissipation doesn't prevent a shear-modified MHD mode from the energy extraction. For the case of helical dissipative flows until now no numerical studies were done. Numerical study of these processes in the swirling flows requires 3-dimensional MHD simulations over large time intervals, which is quite a tough challenge. We expect that the outcome of the ``competition\" between the increase of the energy due to the presence of the shear flow and viscous damping due to the presence of  the dissipation will depend on the time scales of both processes and initial conditions. Relevant numerical studies are planned to be started in the near future.  However, basing on our results about `nonmodal' temporal evolution of Alfv\\'enic perturbations, we can give qualitative picture of the physical appearance of these Alfv\\'en waves in the presence of viscous dissipation for all those three subclasses of flows -- viz: (a) ejectional outflows with no rotation; (b) helical flows with exponentially increasing $|{\\bf k}(t)|$; (c) helical flows with periodically changing $|{\\bf k}(t)|$) -- which were considered.  For an ejectional outflow, with linear temporal growth of the $|{\\bf k}(t)|$, a monochromatic Alfv\\'enic disturbance would undergo {\\it one} act of transient, `burst-like' amplification followed by the viscous decay. For the package, containing harmonics of different $|{\\bf k}(0)|$, one could expect a sequence (`cascade') of their transient amplifications mixed with their viscous damping. This sort of the Alfv\\'en wave, while passing through the flow,  would slowly decay but repeatedly regain its energy via the chain-sequence of amplification events.  For helical flows, where the temporal increase of the  $|{\\bf k}(t)|$ is exponential, the Alfv\\'en waves could be less fortunate. It is true that their energy might grow exponentially, but the increase of the  $|{\\bf k}(t)|$ seems to be a more rapid process than the increase of the total energy of the Alfv\\'en wave (see Fig.6). Therefore, even if initially the Alfv\\'en waves could extract some energy from the background flow, the powerful viscous damping would sooner or later prevail and return the energy back to the flow in the form of heat. It means that in this situation, perhaps, one could not see any high amplitude Alfv\\'en waves, but one would rather witness an efficient heating of the flow due to this {\\it `self-heating'} process: when Alfv\\'en waves serve as energy transmitters, extract a part of the flow's regular kinetic energy and returning the energy back to the flow in the form of the heat.  For helical flows, where the variability of the  $|{\\bf k}(t)|$ is periodic, the situation can be the most exotic. The strength of the viscous damping in such flows is limited, because the value of the $|{\\bf k}(t)|$ is kept finite: it doesn't increase neither algebraically nor exponentially. At the same time, Alfv\\'en waves are modulated by the presence of the flow and there exist separate regions of parametric instability, which implies that under favourable conditions the Alfv\\'en waves propagating through this kind of flow can be unstable. Therefore, flow patterns of this gender might serve as efficient amplifiers of Alfv\\'en waves.  In all these cases the ultimate `fate' of any individual Alfv\\'en wave would depend on the time- and length-scales of their propagation and the flow geometry, as well as, of course, on the actual strength of the viscous dissipation in the flow.  The shear-modified Alfv\\'en waves seem to be the {\\it hybrids} of the usual Alfv\\'en waves and the so called `Kelvin modes' (\\cite{mp77, rcm98}). The latter are nonperiodic, transiently growing modes inherent to the incompressible neutral shear flows. In the absence of the equilibrium magnetic field, the Kelvin Transients are the only modes of collective behaviour exhibited by the flow.  The shear-modified Alfv\\'en waves are formed by combining the features of the usual waves with the transient shear vortices; their ability to extract/give energy from/to the mean flow follows from the latter.  Within the framework of the nonmodal approach we are able to track and study properties of these shear-modified Alfv\\'en waves in the space of wave numbers ($\\bf k$-space). The orientation and magnitude of a wave number vector, being variable in time and governed by Eq. (14), varies because of the local straining of the background flow. The evolution of the ${\\bf k}(t)$ field, in its turn, exercises influence  on the temporal behaviour of physical variables: density and pressure perturbations, components of the magnetic field and the velocity vectors. In terms of the usual physical coordinates (space and time) these variables are periodic in each space coordinate. However,  these modes are `non-normal' (\\cite{bf92}) because both wavenumbers and amplitudes associated with each disturbance are functions of time, implying that the solutions are {\\it not} separable in space and time. The latter, nonexponential time-dependence of the solutions is disclosed through the nonmodal approach and it helps to ``smoke-out\" new classes of modes hardly accessible through traditional normal mode analysis. In order to have complete real-space physical picture of the evolution of these Alfv\\'enic perturbations we need, therefore, to develop phenomenological approach and try to understand the dynamics in terms of vorticity dynamics (\\cite{cc86}). Alternatively one might approach the problem through direct numerical simulations, generalizing the very first attempts of numerical studies of much simpler MHD velocity patterns (\\cite{bprr01}).  Turning our attention back to astrophysics, we might argue that the role of the shear-modified hybrid Alfv\\'en waves could be quite important in a number of astrophysical situations. Let us first assess the relevance these waves could have to the {\\it solar plasma flows}. It is known that higher layers of the solar chromosphere consist of long, vertical columns ({\\it spicules}) rising above the general background (\\cite{a86}) and well-visible at the extreme limb of the Sun, where the overall structure of the chromosphere resembles a ``burning prairie.\"  The spicules rise vertically, like slender magnetic flux tubes (filaments) out of the network beginning in the chromosphere and threading through the solar transition region into the low corona. Their average lifetime is 10-15 minutes, characteristic length scale $10.000 km$, and characteristic diameter $700-1000 km$. Often the spicules appear to have an ``Eifel tower\" shape with vertical flow velocities in the range $20-25 km/s$. There is also some evidence for plasma rotation in spicules (\\cite{ks99}). It means that the plasma motion within spicules may easily be of helical nature. The recent discovery of swirling macrospicules, called {\\it solar tornados} (\\cite{pm98}) provides additional evidence in this direction.  If we imagine flow patterns, like spicules or solar tornadoes, with moderate storage of total (quasi)equilibrium energy it may happen that transiently amplified Alfv\\'en waves, bred within these structures, could efficiently `suck out' much of their energy and lead to eventual exhaustion and disappearance of the flow patterns. The shear-induced cascade of Alfv\\'en waves could take away much of the available energy of the ``parent\" flow. The observational consequences may be dramatic; we could imagine individual spicules acting as short-lived plasma cannons firing away heavy blasts of large-amplitude Alfv\\'en waves and disappearing after the wave package has drained most of the energy from the ``parent\" flow.  It is quite fascinating to surmise that cascade of transient amplifications of Alfv\\'en waves within spicules, together with other kinds of shear-induced wave processes (like, e.g., ``MHD wave oscillations\" \\cite{rpm00}) may account for the quasisteady appearance of spicules in the solar atmosphere. Taking the ``burning prairie\" metaphor a little further,  we may speculate that shear induced processes, ``burning\"  individual ``grass blades\", could be responsible for the overall conflagration of the spicule network which, in turn, provides an efficient mechanism for the transfer of energy from the quasisteady motions of solar plasmas into the flux of large-amplitude Alfv\\'en waves. These waves are able to penetrate through the solar transition region and reach the lower layers of the solar corona. Here the processes of the shear-induced ``MHD wave oscillations\" (\\cite{rpm00}) may contribute to the appearance of compressible MHD modes (slow and/or fast magnetosonic waves) which, in turn, are effectively damped giving back their energy to the solar plasma particles, accelerating them and providing initial launch of the solar wind flows. This highly speculative picture must be  carefully developed and tested in order to make it convincing. We are in the process of developing a numerical model for an individual helical solar flow pattern, and hope to see the above-described processes in  real physical space through  numerical simulations.  The parametric Alfv\\'en instability  may be effective in flows, where conditions are favorable for the correlation between the rate of rotation and stretching of the flow lines (which determines the value of the characteristic frequency $W$) and the strength of the equilibrium magnetic field (which determines the magnitude of the Alfv\\'en frequency ${\\omega}_{\\rm A}$). Such a correlation may easily happen within solar jet flows, where both the ejectional and rotational modes of motion are present. Any time such correlation takes place it will lead to a burst-like generation of high-amplitude Alfv\\'en waves with corresponding ``exhaustion\" (and, maybe, even disappearance) of the ``parent\" helical flow pattern. Since the modern apparatus used on the last generation solar satellites enables one to measure both the strength of magnetic fields and the rate of rotation, one can assume that it will be fairly realistic to verify whether the appearance of high-amplitude Alfv\\'en waves is really related to shear-induced resonant parametric instability of helical flows.  We think that shear-induced instabilities of this nature may be also present in {\\it accretion-ejection flows}.  The differential rotation parameter $n$ controls the ${\\bf k}(t)$  dynamics which, in turn, specifies the nature of fluctuations in a given flow. When $n<1$ (including the rigid rotation case) $\\Gamma$ is imaginary and ${\\bf k}(t)$ is periodic. In such cases (e.g., innermost regions of galactic gaseous disk, believed to have almost constant rotation rates) the system may sustain parametrically unstable Alfv\\'en waves. While when $n>1$ (including the Keplerian rotation regime), $\\Gamma$ is real and makes the time behavior of ${\\bf k}(t)$ exponential. This regime can be realized in different kinds of quasi Keplerian accretion disks. Note that the characteristic amplification time scale ${\\cal T}{\\simeq}{\\Omega}^{-1}$ is of the order of  the inverse period of rotation, being at least of the same order as the ``Magnetorotational instability\" growth rate (\\cite{bh98}). Since the latter instability is thought to be the strongest one, accounting for the turbulence in accretion disks, we could surmise that the shear-induced instability discussed  in this paper could also play a role in the onset of  turbulence in accretion-ejection flows.  Finally, considering our results in the context of galactic gaseous disks and their large-scale magnetic fields, we might envisage a possible relation between the shear-induced amplification of Alfv\\'en waves and standard dynamo action in swirling flows. Within galactic disks algebraic and/or usual instabilities of the nonmodal origin might provide the important mechanism for the initial amplification of the magnetic energy and this process could compete with the omnipresent Ohmic diffusion. This is a challenging task to see what is  the role of nonmodal, shear-induced processes in the framework of the general theory of dynamo generation of galactic magnetic fields."
    },
    "0212/astro-ph0212304_arXiv.txt": {
        "abstract": "Energetic outflows appear to occur in conjunction with active mass accretion onto supermassive black holes.  These outflows are most readily observed in the $\\sim10\\%$ of quasars with broad absorption lines, where the observer's line of sight passes through the wind. Until fairly recently, the paucity of X-ray data from these objects was notable, but now sensitive hard-band missions such as \\chandra\\ and \\xmm\\ are routinely detecting broad absorption line quasars. The X-ray regime offers qualitatively new information for the understanding of these objects, and these new results must be taken into account in theoretical modeling of quasar winds. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/astro-ph0212148_arXiv.txt": {
        "abstract": " ",
        "introduction": "The recent cosmological observations show that the universe is in an accelerating epoch.  These observations, together with CMB data, require a dark energy density of $\\Omega_{de} \\simeq 0.7$ and negative equation of state $w_{de}=p_{de}/\\rho_{de}<0$ \\ci{CMBR,SN1a}. The data suggest that $w_{de}<-2/3$ and the best fit models have $w_{de}$ very close to minus one   including the regions around $w_{de}=-1$. Even more, the existing data do not ruled out a $w_{de}<-1$ \\ci{SN1a}. Causality arguments seem to require $w>-1$ but this is only the case for constant $w$ and in fact it is possible to have $|w|>1$ without violating causality \\ci{ax.vuc}. The possibility of having $w<-1$ was raised by \\ci{cald}-\\ci{mcin} where they studied a phantom energy density with $\\rho_{ph}>0$ and $w_{ph}<-1$. From a particle point of view it is unclear how such field can be obtained and it was suggested to have a scalar field with a negative kinetic energy \\ci{cald}, \\ci{phantom}. In this work, we will concentrate on scalar fields with canonical kinetic terms ($E_k\\geq0$) that lead dynamically to regions with $|w|>1$.  The cosmology of negative potential has become increasingly interesting since many particle physics models and string theory predict anti-de-Sitter \"ADS\" spaces. The cyclic universe \\ci{cyc} and the ekpyrotic \\ci{ekp} universe models use a negative potential and describe  a universe that goes from a contracting phase (leading to a big crunch) to an accelerating universe. The main idea is to give an alternative theory to the big bang singularity. However, it is not yet clear the process of the transition between contracting and accelerating phases.    Scalar fields are widely obtained in particle physics, and they can be either fundamental fields or compose fields (fermion condensates) \\ci{chris1,ax.mod}. They can be  used to describe the cosmological constant as quintessence \\ci{tracker} and in all these models the scalar fields have a positive semi-definite potential $V$ \\ci{generic, mio.scalar}. However, negative potential are widely predicted by string theory or  by supergravity (sugra) models. The sugra potential is \\be V=e^K\\le[F_i(K^{-1})^i_jF^j-3|W|^2\\ri] \\ee with $K(\\phi,\\bar\\phi)$ the Kahler potential, $W(\\phi)$ the superpotential and $F_i=K_iW+W_i,\\;W_i\\equiv dW/d\\phi_i$. It is clear that the general vacua of sugra  has a negative potential and a lot of model building (in many cases not quite self consistent) has to be done to get a potential $V\\geq 0$.   Negative potentials were studied in \\ci{linde} but in the absence of a background density and in a model dependent analysis. Here, we will study the cosmology of negative potentials in the presence of a background energy density, we do a model independent analysis of the evolution of the scalar field and we give some specific examples. We will show that these models lead to an equation of state with $|\\wp|>1$ in a natural dynamical way. The stage where  $\\wp < -1$  has a negative energy density $\\rho_\\phi <0$ but a positive $\\rp+p_\\phi=2E_k\\geq 0$. These models  lead inevitable to a collapsing universe and to a would be \"big crunch\" contrary to phantom fields. Indeed, phantom physics suggest that the final stage of the universe has a divergence of the scale factor  at finite time ($a (t) \\rightarrow \\infty)$ \\ci{mcin} called a \"big smash\" while a \"big crunch\" ends in a singularity $a(t) \\rightarrow 0$, $\\rho\\rightarrow \\infty$.   We do not worry about a big crunch as we do not worry about a  big bang since we are working with an effective theory and we would expect that a fundamental theory like strings could have a graceful exit to the problem of singularities. The behavior of the universe approaching the  big crunch singularity is just as bad  as the big bang singularity since the universe has, in general, a time inversion symmetry $t_0+t\\rightarrow t_0-t$ around the point where the Hubble constant is zero  $H(t_0)=0$. So, whatever the solution to the big bang singularity will also solve the big crunch singularity. The evolution of the universe outside the possible singularities  may well be described by the effective field theory and general relativity.   An accelerating universe requires a dominating energy density with $w_i<-1/3$, (or more exact $\\Omega_i w_i<-1/3,\\;\\Omega_i=\\rho_i/\\rho_{Tot}$), and therefore we could naively think that a $w_i<-1$ will increase the acceleration of the universe. However, a canonical normalized scalar field with $\\wp<-1$ decelerates the universe because in this region the energy density is negative, as we will show later, and its contribution to the acceleration of the scale factor  given by \\be \\ddot{a}\\sim - (\\rho_\\phi+3p_\\phi)=-\\rho_\\phi(1+3\\wp) \\ee is negative, thus decelerating the expansion of the universe.    In particle physics the prediction of many scalar fields is quite common. It is certainly more natural to assume the existence of more than one scalar field which will, in general, have different potentials and different cosmological evolutions. So, can quintessence  be the sum of two or more different kinds of scalar fields? The answer is clearly yes. There is no theoretical or observational argument against it.  The interesting point is that negative potentials allow, al least in principle, to explain the cosmological scale in terms of a high energy scale (e.g. electroweak  or susy breaking scale) since the contribution of a negative potential will cancel that of a positive potential rendering a finite small positive  cosmological or quintessence energy \\ci{vilenkin}. Of course, there is a fine tuning problem such that the difference of two quantities at the  high energy scale  gives something of the order of today's energy $\\Lm_c\\sim 10^{-12} GeV$ and that such a cancellation must take place during nucleosynthesis and today.   However,   the evolution of the scalar field with negative potential will necessarily   give a   region with a small cosmological constant, since for $\\rho_\\phi$ negative we have $|\\rho_\\phi|<\\rho_b$ regardless of the functional form of the potential $V$ and the fact that $\\wp<-1$ which makes $|\\dot\\rho_\\phi|$ grow faster than $\\rho_b$. The evolution is such that $\\rho_b+\\rho_\\phi\\rightarrow 0$ for any potential $V$ or energy density $\\rho_b$. So it is an attractor solution.  During all the time that $\\rho_\\phi<0$ we have $|\\rho_\\phi| < \\rho_b$, i.e. it will get smaller than the sum of all other energy densities for a long period of time and regardless of the functional form of the potential $V$. ",
        "conclusions": ""
    },
    "0212/astro-ph0212181_arXiv.txt": {
        "abstract": "We assess the possibility of determining the Hubble constant $H_0$ by measuring time delays between multiple images of supernovae gravitationally lensed by rich clusters of galaxies and combining these delay measurements with \\mbox{detailed} cluster-potential models based on other lensing constraints. Such a lensing \\mbox{determination} of $H_0$ would be complementary to those obtained from galaxy-QSO lensing studies, and could potentially be better calibrated. We show that relatively low-redshift ($z \\sim 0.2$), significantly elliptical clusters have appreciable lensing cross sections for observable image pairings with tractable time delays on the order of a few years despite large lensing mass scales. We find that a targeted search for such image pairs would be a significant undertaking for current observatories, but that it would be appropriate for a facility such as the proposed Large-aperture Synoptic Survey Telescope. ",
        "introduction": "The best measurements of the Hubble constant ($H_0$) based on the Cepheid-calibrated distance ladder and on the cosmic microwave background agree impressively with one another: $H_0 = 72 \\pm 8 \\, \\mathrm{km} \\, \\mathrm{s}^{-1} \\, \\mathrm{Mpc}^{-1}$ and $H_0 = 72 \\pm 5 \\, \\mathrm{km} \\, \\mathrm{s}^{-1} \\, \\mathrm{Mpc}^{-1}$, respectively \\citep{fre01, spe03}.  However, these values are inconsistent with much lower values of $H_0$ based upon measurements of gravitational lens time delays and the assumption that lensing galaxies have extended dark halos \\citep{koc02a, koc02b, koc03}. In this paper we explore the possibility of determining $H_0$ by monitoring known strong-lensing galaxy clusters and measuring time delays between multiple images of high-redshift supernovae (SNe).  The number of known giant-arc clusters is encouraging. \\citet{wu98} present a summary of 38 strongly lensing clusters from the literature, containing a total of 48 giant arcs and arclets (see also \\citet{wnb99}).  More recently, \\citet{lup99} have reported the discovery of 8 more giant-arc clusters, \\citet{gye02} 6 more, and \\citet{zg03} 3 more. Recent observations of the galaxy cluster Abell 1689 with the Advanced Camera for Surveys (ACS) on the {\\sl Hubble Space Telescope} have revealed at least 30 new multiply-imaged background sources, with an average of $\\sim$3 images each (D. Coe, 9 January 2003 presentation ``Deep ACS and Keck Observations of A1689'', in conference ``Gravitational Lensing: a Unique Tool for Cosmology'', Aussois, Savoie, France). All these image systems (together with previously known strongly lensed features) provide constraints on the projected gravitational potential of the cluster---many more constraints than a quadruple-image quasar lens system places on its lensing galaxy. In addition, since multiple strongly lensed sources behind a single galaxy cluster in general lie at different redshifts, the cluster potential map will not suffer the ``mass sheet degeneracy'' that afflicts galaxy-quasar lens systems \\citep{sah00}. Although A1689 is perhaps the strongest cluster lens in the sky, we may assume that other known strong-lensing clusters will likewise exhibit many new multiple-image systems when observed with ACS and later instruments, and detailed cluster modeling like that of \\citet{asw98a, asw98b} and \\citet{k95, k96} but with increased accuracy will become possible.  Thus although the structure of galaxy clusters is in general more complex than that of individual galaxies, the projected potentials of clusters may in the near future be more tightly constrained by lensing than those of individual galaxies, and cluster-lens time delay would therefore be more precisely calibrated. (See \\citet{sch00} for a general discussion of the determination of $H_0$ from gravitational-lens time delays.)  This work is distinguished from previous studies by its quantitative focus on the possibility of measuring SN-image time delays in cluster-scale lens systems, and combining them with additional strong-lensing constraints to determine $H_0$. The original proposition to measure $H_0$ using gravitational lens time delays envisioned distant supernovae (SNe) as the variable sources to be lensed by intervening galaxies \\citep{ref64}; quasars and their variability had yet to be discovered. \\citet{kp88} examined the possibility of observing SNe exploding within giant-arc source galaxies. They made a detailed analysis for a SN very near to and inside a caustic ``cusp'' and showed that for such configurations the time delay between multiple images could be as short as days to weeks, although they did not explicitly consider the possibility of making an $H_0$ determination.  \\citet{gir92} reported the detection of variability in two separate arclets of the cluster Cl 0302+1658, consistent with a single source and a time delay of 1 to 2 years. Various papers have examined magnification and time delay effects for SNe in field surveys lensed by intervening halos across a spectrum of mass scales \\citep{lsw88, mf98, df99, mfp00, pm00, holz01, goo02, ost03, ok03}, and several have explicitly considered the possibility of constraining $H_0$ using SN time delays.  However, all these works have modeled the lenses as circular, and in this paper we show that circular models seriously under-predict the capacity for cluster-scale lenses to produce multiple images with reasonably short time delays (as one might anticipate based on the results of \\citet{kp88} for sources very near to the caustic cusp of a lens). Other studies have considered targeted searches for lensed SNe in galaxy cluster fields \\citep{mr97, kb98, srs00, sul00, gms02, gg02}, but their quantitative focus has been on pure SN detection rates and not on the prospects for measuring time delays that we quantify here. With the exception of Saini et al.\\ (2000), these works have also approximated clusters with circularly symmetric lens models.  This paper is organized as follows. Section~\\ref{theory} presents a brief discussion of relevant gravitational lens theory and defines the particular lensing cross section that is central to our calculations. In Section~\\ref{cross_section}, we examine the lensing behavior of several cluster models in the context of observational constraints on minimum magnification and maximum time delay. Section~\\ref{accuracy} discusses the accuracy in the determination of $H_0$ that we might hope to obtain from individual cluster time delays, both for a simplified model and for realistic models. Section~\\ref{detect_sec} presents a calculation of detection rates for observable SN image pairings with acceptable time delays in the strong-lensing region of a specific cluster. Our approach in this section is similar to that of \\citet{sul00} and \\citet{gg02} (and differs from that of Saini et al. (2000)) in that we do not directly consider the possibility of SNe within known lensed galaxies, but rather attempt a prediction of SN detection rates in a targeted cluster field based on based on an assumed cluster mass model and cosmic supernova rate density function. (\\citet{lan02} argue that above redshifts of $\\sim$1.5--2, even the Hubble Deep Field observations are insensitive to the surface brightness of most rest-frame ultraviolet emission.) Finally, in Section~\\ref{outlook} we discuss the observational outlook in view of the calculations of Section~\\ref{detect_sec}. We find that a measurement of $H_0$ based on cluster time delays would be challenging but not impossible with today's telescopes and instruments, and that it would be an appropriate project for a large telescope operating in a dedicated survey mode such as the proposed Large-aperture Synoptic Survey Telescope ({\\sl LSST}) \\citep{ty02}.  \\enlargethispage*{1000pt} Throughout this paper, we assume a flat, vacuum-dominated cold dark matter ($\\Lambda$CDM) cosmology with matter and vacuum density parameters $\\Omega_{\\mathrm{M}} = 0.3$, $\\Omega_{\\Lambda} = 0.7$. For the Hubble constant, we take $H_0 = 70 \\, h_{70} \\, \\mathrm{km} \\, \\mathrm{s}^{-1} \\, \\mathrm{Mpc}^{-1}$ with $h_{70} = 1$. \\pagebreak ",
        "conclusions": ""
    },
    "0212/astro-ph0212462_arXiv.txt": {
        "abstract": "Narrow-band $HST$ WFPC2 images reveal a bow-shock-like halo around the \\hii\\ region N30B toward the B[e] supergiant Hen S22 located within the larger DEM\\,L\\,106 nebula in the Large Magellanic Cloud. High-dispersion spectra of N30B show a narrow \\ha\\ emission component from the ionized gas; the velocity variations indicate a gas flow of $-$5 to $-$10 \\kms\\ in the vicinity of the \\hii\\ regions, which is resultant from interactions with Hen S22's stellar wind and responsible for the bow-shock morphology. Spectra of N30B's halo show broad \\ha\\ profiles extending over $>1000$ \\kms, similar to that of Hen S22, indicating that the halo is a reflection nebula of Hen S22. Broad-band morphologies of N30B's halo are also consistent with the reflection nebula interpretation. We use dust-scattering properties and the observed brightnesses of the reflection nebula and Hen S22 to constrain the reflection geometry. The reflected stellar \\ha\\ emission and absorption vary across the reflection nebula as a result of viewing S22's anisotropic wind across varying angles.  This reflection nebula, together with the edge-on orientation of Hen S22's disk, provides an invaluable opportunity to study the disk and polar winds of a B[e] supergiant. ",
        "introduction": "A {\\it Hubble Space Telescope (HST)} WFPC2 \\ha\\ image of the \\hii\\ region \\dem\\ \\citep{DEM} in the Large Magellanic Cloud (LMC) revealed a pair of bright, compact \\hii\\ regions embedded within a large shell nebula \\citep{chenetal00}.  \\dem\\ has been previously cataloged by \\citet{H56} as N30, and the two compact \\hii\\ regions were identified as an \\hii\\ knot and designated as N30B.  As shown in Figure~1, while the main body of N30B is unremarkable, it is surrounded by a halo with sharp edges on the northern rim facing the luminous B[e] supergiant Hen S22 \\citep[= HD\\,34664;][]{H56}.  B[e] supergiants are known to possess stellar winds \\citep{Zetal96}.  The nebular morphology and relative location of Hen S22 and N30B thus appear to suggest that the halo of N30B is a bow shock produced by dynamical interactions with the wind of Hen S22.  A two-component stellar wind model has been proposed for B[e] supergiants: (1) a cool, slow wind along a dusty equatorial disk which gives rise to narrow, low-excitation emission lines of, e.g., \\ion{Fe}{2}, [\\ion{Fe}{2}], and [\\ion{O}{1}], and (2) a hot, fast polar wind that is responsible for the broad, high-excitation absorption lines of, e.g., \\ion{C}{4}, \\ion{Si}{4}, and \\ion{N}{5} \\citep{Zetal85,Zetal86}. $IUE$ spectra of Hen S22 in the 1800--3200 \\AA\\ range are dominated by P Cygni type lines of \\ion{Fe}{2}, from which it is derived that the equatorial disk is viewed nearly edge-on and the terminal wind velocity along the disk is $\\sim80$ \\kms \\citep{Zetal96}; while $IUE$ spectra in the 1150--1950 \\AA\\ range show blue-shifted absorption of \\ion{Si}{4} and \\ion{C}{4} indicating that the polar wind of Hen S22 has an expansion velocity $>$1000 \\kms\\ \\citep{Betal83,Zetal86}. N30B is projected at 20$''$--30$''$ \\citep[or 5--7.5 pc for a distance of 50 kpc,][]{F99} from Hen S22; its 24$''$ spatial extent (Figure 1) has a good chance to be partially located above the disk and interact with the polar wind of Hen S22.  To determine the nature and physical properties of the bow-shock-like halo around N30B, we have obtained high-dispersion, long-slit echelle spectra of N30B, its bow-shock-like halo, its surrounding interstellar medium, and Hen S22.  We find that the bow-shock-like halo shows surprisingly broad \\ha\\ line profiles similar to that of Hen S22, suggesting that it is a reflection nebula. Furthermore, the spectral features vary from the east side to the west side of N30B, indicating that {\\it N30B indeed sees the wind of Hen S22 along a range of viewing angles}. Recently archived $HST$ WFPC2 images of N30B obtained by the Hubble Heritage Project provide additional support to the interpretation of a reflection nebula for the halo of N30B. In this paper, we report these observations, and discuss the physical properties of N30B, the nature of its halo, and the implication of our results on the disk and polar winds of the B[e] supergiant Hen S22. ",
        "conclusions": "The two small \\hii\\ regions in N30B are aligned roughly along the east-west direction, thus we call these two regions N30B-E and N30B-W, respectively. In the \\ha\\ image, apart from the bow-shock-like halo, the bright central regions of these two \\hii\\ regions have simple morphologies expected for Str\\\"omgren spheres. N30B-E is about 7$''$ (1.7 pc) in diameter and shows a bright spot at its northern edge, while N30B-W is 6$''$ (1.45 pc) across and has a rather uniform surface brightness.   \\subsection{Densities and Ionization of the \\hii\\ Regions in N30B}  The $HST$ WFPC2 \\ha\\ image, with stellar emission excised, can be used to estimate the \\ha\\ luminosities and required ionizing fluxes of the \\hii\\ regions in N30B. We have made two measurements which differ in their assumption as to whether the halo is a reflection nebula or part of the \\hii\\ regions. If the halo around the \\hii\\ regions is an envelope reflecting the light from Hen S22, a background component similar to the halo should be subtracted. The resultant integrated \\ha\\ fluxes for the 7$''$ and 6$''$ Str\\\"omgren spheres of N30B-E and N30B-W are $(5.0\\pm0.5)\\times10^{-13}$ ergs~cm$^{-2}$~s$^{-1}$ and $(3.3\\pm0.4)\\times10^{-13}$ ergs~cm$^{-2}$~s$^{-1}$, respectively. These fluxes have been corrected for the filter transmission and \\nii\\ contamination. If the halo is ionized, a low background exterior to N30B should be used and the integrated fluxes of the \\hii\\ regions including the halo would be about 80\\% larger; however, as we show in \\S3.3 this is not the case, so these large fluxes will not be used in further discussions. Adopting the extinction of D106-10 in N30B-W, $E(B-V)$ = 0.14 \\citep{oey96a}, and a distance of 50 kpc, we derive \\ha\\ luminosities of $(2.0\\pm0.2)\\times10^{35}$ ergs~s$^{-1}$ and $(1.4\\pm0.2)\\times10^{35}$ ergs~s$^{-1}$ and required ionizing fluxes of $(1.7\\pm0.2)\\times10^{47}$ photons~s$^{-1}$ and $(1.1\\pm0.2)\\times10^{47}$ photons~s$^{-1}$ for the Str\\\"omgren spheres of N30B-E and N30B-W, respectively. Using the average surface brightness and adopting an emitting path length similar to the diameter of the \\hii\\ region, we have further derived the rms densities of N30B-E and N30B-W, both $\\sim$75 H-atoms cm$^{-3}$.  To examine the stellar content and to estimate the available ionizing fluxes in N30B, we have carried out photometry for stars brighter than $V$ = 20, using the WFPC2 $B$ and $V$ images. Such bright stars are expected to be of early types, so it is possible to apply constant zeropoints to transform the WFPC2 $B$ and $V$ to Johnson $B$ and $V$. The color-magnitude diagram of Johnson $V$ versus $B-V$ is presented in Figure 3. Using the colors and magnitudes and assuming that these stars are still on the main sequence, we can deredden the stars and estimate their spectral types. This is easily done by projecting the position of a star in the color-magnitude diagram along the dereddening direction back to the main sequence. The photometric measurements and estimated spectral types are given in Table 2. Most of these stars have also been studied by \\citet{oey96b} using ground-based observations. The differences between her results and the new WFPC2 results are also given in Table 2 for comparison. Excellent agreements in photometry are found for stars in isolation, but the ground-based measurements for faint stars in crowded regions may have errors as large as $\\Delta V \\sim$ 1 mag.  In N30B and its vicinity, stars brighter than $V$ = 20 are all B stars. Ionizing fluxes from early-type stars have been given by \\citet{p73} down to B3 and \\citet{SdK97} down to B0.5. The ionizing fluxes of B0\\,V and B0.5\\,V stars given by the latter are significantly higher than those of the former. N30B-E contains only one early B star, star \\#13 in Table 2; the ionizing flux of this B2.5\\,V star is $<10^{45}$ photons s$^{-1}$ \\citep{p73}, two orders of magnitude lower than that required to ionize N30B-E. It is possible that the B[e] supergiant Hen S22, of spectral type B0--B0.5 \\citep{Zetal86}, plays a significant role in the photoionization of N30B-E. N30B-W contains the bright star \\#6, which is saturated in the WFPC2 images, but has been previously classified spectroscopically as a B1.5\\,V star \\citep[= D106-10;][]{oey96a}. The ionizing flux of a B1.5\\,V star given by \\citet{p73} is 2$\\times$10$^{45}$ photons s$^{-1}$, which is also about two orders of magnitude lower than the required. However, extrapolating from the ionizing fluxes of B0\\,V and B0.5\\,V stars given by \\citet{SdK97}, we find that the ionizing flux of the B1.5\\,V star \\#6 (D106-10) is sufficient to ionize N30B-W.   \\subsection{Relationship between N30B and the Giant Shell DEM\\,L\\,106}  The relationship between the \\hii\\ regions in N30B and the surrounding giant shell DEM\\,L\\,106 may be assessed from their kinematic properties. An example of our echelle observations is illustrated in Figures 4a and 4b in two stretches to show the \\ha\\ and \\nii\\ $\\lambda$6548 lines from the E-W slit centered near the star \\#6 in N30B-W. The vertical axis is spatial with east on top, and the horizontal axis is spectral with wavelength increasing to the right. For convenience, we have converted wavelengths to heliocentric velocities of the \\ha\\ line and marked them along the horizontal axis in Figure 4. (All radial velocities in this paper are heliocentric.) The narrowest lines with uniform surface brightness along the slit are telluric \\ha\\ and OH lines, located near 0, $-$425, and +660 \\kms. The nebular \\ha\\ and \\nii\\ $\\lambda$6548 lines can be found at +300 and $-$380 \\kms, respectively. The \\ha\\ line in Figures 4a and 4b shows a bright, narrow component and a faint, broad component. The spatial extent of the broad component coincides with that of N30B's halo; this component will be discussed in \\S3.3. The narrow \\ha\\ component arises from ionized gas. The two \\ha\\ emission peaks near the center of Figure 4a correspond to the \\hii\\ regions N30B-E and N30B-W, while the fainter \\ha\\ emission extending throughout the entire slit length corresponds to the giant shell in DEM\\,L\\,106. No line splitting is detected in \\ha\\ or \\nii\\ for either N30B or DEM\\,L\\,106.  To analyze the kinematics of ionized gas quantitatively, we have extracted velocity profiles of the \\ha\\ and \\nii\\ lines along each slit, made Gaussian fits, and plotted the results in Figure 5. The fitted velocities are plotted along the slit positions, with \\ha\\ in filled symbols and \\nii\\ in open symbols; the fitted FWHM of the \\ha\\ line is plotted as the ``error bar\" at each position, except where data points are densely packed. The observed FWHM of the \\ha\\ line is the quadratic sum of the instrumental FWHM (13.5 \\kms), the thermal FWHM (21.4 \\kms\\ at a temperature of 10$^4$ K), and the turbulent FWHM. Thus the observed \\ha\\ FWHM of $\\sim$26 \\kms\\ for the \\hii\\ regions in N30B implies a turbulent FWHM of $\\sim$6 \\kms. This small internal turbulence is consistent with the expectation for \\hii\\ regions ionized by early B-type stars, because they lack strong fast stellar winds. The giant shell in DEM\\,L\\,106 has a larger observed FWHM in \\ha, 35--40 \\kms, implying a turbulent FWHM of 24--31 \\kms. The lack of line splitting suggest that this large turbulent FWHM is dominated by turbulence instead of a systematic shell expansion.  Further kinematic information is provided by the velocity variations shown in Figure 5. The average heliocentric velocity of DEM\\,L\\,106 is $\\sim$ 300 \\kms. Gradual velocity excursions of up to $\\pm$10 \\kms\\ from the average velocity are seen. Along the NS slit centered on Hen S22, a systematic blue-shift of $-$10 \\kms\\ is seen in the \\nii\\ $\\lambda$6583 line. (The nebular \\ha\\ line is too contaminated by the stellar emission to be measured near the star.) These blue-shifts provide the only evidence that the internal motion of DEM\\,L\\,106 does contain a component of systematic expansion at a level of $\\sim$10 \\kms. The kinematics of DEM\\,L\\,106 show that this giant shell, formed by Hen S22 and the other massive stars in the OB association LH38 \\citep{LH70}, is by no means a simple expanding shell with a large central cavity. The interior of DEM\\,L\\,106 may still contain a significant amount of gas in the form of dense cloudlets, and the dynamical interactions between the fast stellar winds and the cloudlets contribute to the large turbulent FWHM observed in DEM\\,L\\,106.  The systemic velocities of the \\hii\\ regions in N30B are similar to that of DEM\\,L\\,106, suggesting that N30B is physically associated with DEM\\,L\\,106. The gas in the vicinity of N30B and especially between its two \\hii\\ regions shows blue-shifts of $-$5 to $-$10 \\kms\\ from the systemic velocity (see Figure 5). Low-velocity flows can be produced by an evaporation off the ionization fronts in \\hii\\ regions or dynamic interactions between stellar winds and \\hii\\ regions. In an evaporation flow, the gas projected outside the \\hii\\ disk would have motions perpendicular to the line of sight and show no offsets from the systemic velocity. In a dynamic interaction scenario, if the stellar wind source is on the far side of an \\hii\\ region, the ablation of the \\hii\\ region by the stellar wind would produce a gas flow toward us. The ubiquitous blue-shifts observed in the vicinity of N30B indicate that the dynamic interaction between the stellar wind of Hen S22 and the \\hii\\ regions in N30B is responsible for these low-velocity gas flow. N30B is probably one of the dense clouds within DEM\\,L\\,106; its visibility is attributed to the early-B stars formed in its interior.  \\subsection{Spectra and Physical Nature of the Halo of N30B}  We initially speculated that the bow-shock-like halo around N30B was produced by interactions between the fast polar wind of Hen S22 and the \\hii\\ regions, and expected to see corresponding supersonic kinematic features within the spatial extent of the halo. The \\ha+\\nii\\ echellogram of N30B in Figure 4 shows that the only kinematic feature that matches the spatial extent of the halo is a prominent, broad \\ha\\ emission component. The continuous distribution of this broad component over 18$''$, from the eastern edge of N30B-E's halo to the western edge of N30B-W's halo, indicates that the halo envelopes the \\hii\\ regions of N30B. The broad \\ha\\ emission is centered near 300 \\kms\\ and detected over a velocity range greater than 1000 \\kms. Such high velocities obviously cannot be accelerated by the wind of Hen S22; furthermore, such strong shocks would have produced more dramatic dynamic effects, such as X-ray emission, which is not seen \\citep{chuetal95}.  Large velocity widths ($>$ 1000 \\kms) have been observed in other \\hii\\ regions, for example, the Carina Nebula in our Galaxy and the giant \\hii\\ region NGC\\,5471 in M101. The broad hydrogen Balmer line emission in the Car II region in the Carina Nebula has been attributed to the dust-scattered Balmer emission of $\\eta$ Carinae and is detected at as far as 5 pc from $\\eta$ Car \\citep{LM86}. The broad \\ha\\ and \\nii\\ lines in NGC\\,5471 have been suggested to be powered by supernova explosions and fast stellar winds \\citep{cvc90}. Clearly, N30B cannot be associated with supernova explosions, since no X-ray emission has been detected from this region \\citep{chuetal95}. Furthermore, N30B does not show broad \\nii\\ lines as are seen in NGC\\,5471, and therefore it should not have the same physical nature as the broad-line region in NGC\\,5471. We will thus focus on whether the broad \\ha\\ line from the halo of N30B may be dust-scattered \\ha\\ emission of the B[e] supergiant Hen S22.  To compare the broad \\ha\\ emission from the halo of N30B to that of Hen S22, we display the echellogram and \\ha\\ line profiles of Hen S22 in Figure 4c--d, below the echellogram of N30B. While they show qualitatively similar profiles with an absorption component blue-shifted from the centroid of the broad \\ha\\ emission, the widths and depths of the absorption are different. A similar behavior is observed in the reflection nebula Car II near $\\eta$ Car. \\citet{LM86} extracted \\ha\\ line profiles from six positions in Car II; these broad profiles (in their Figure 3) all show an absorption component on the blue side of the nebular line, but the absorption strength decreases radially from $\\eta$ Car. \\citet{LM86} suggested that the spatial variations in the dust-scattered \\ha\\ line in Car II reflected the temporal variability of the \\ha\\ profile of $\\eta$ Car \\citep{Retal84}. However, after monitoring the \\ha\\ profiles of $\\eta$ Car and Car II from 1985 to 1997, \\citet{Betal98} show that the the absorption feature in the reflected \\ha\\ line is spatially variable across Car II but temporally invariable over the 12-yr span, and conclude that the reflection nebula is located within a 46$^\\circ$-wide cone of light from $\\eta$ Car along the axis of its bipolar Homunculus nebula. We suggest that the broad \\ha\\ line from N30B's halo is also dust-scattered \\ha\\ emission from a nearby \\ha-emitting star, Hen S22 in this case. The position-dependence of the broad \\ha\\ line in N30B is also caused by different viewing angles, instead of a temporal variability of Hen S22, as discussed further in \\S3.5.  The reflection nebula around N30B is not unique in the vicinity of Hen S22.  Our echelle observations also detected broad \\ha\\ line profiles in the nebulosity at 20$''$ east of N30B, at a much fainter level. The bow-shock-like morphology of the reflection nebula around N30B is also not unique among reflection nebulae, as a similar nebula has been observed around the B2\\,V star LSS\\,3027 in our Galaxy \\citep{chu83}. In both cases, the bow-shock-like structure curves with the apex toward an OB association, LH\\,38 for N30B \\citep{oey96a} and Stock 16 for LSS\\,3027 \\citep{chu83}. The relative nebular and stellar positions are highly suggestive that the stellar radiation and winds from the OB association are responsible for the bow-shock-like morphology of the reflection nebula. For N30B, the dynamical interaction is further supported by the small blue-shift, $-$5 to $-$10 \\kms, of the ionized gas in the vicinity of the two \\hii\\ regions. This flow velocity, $-$5 to $-$10 \\kms, is subsonic for 10$^4$ K ionized gas but supersonic for cold neutral gas. Since the reflective envelope of N30B contains dust and mostly neutral gas, flow velocities of 5--10 \\kms\\ would be supersonic; thus the morphology of N30B's halo may indeed be produced by a bow shock.   \\subsection{Geometry of the Reflection}  The nature of N30B's halo as a reflection nebula is supported by the WFPC2 $B$, $V$, and $I$ images in Figure 2. These three broad-band images show similar morphologies among one another, with almost uniform surface brightness across N30B. Hardly any enhancement is seen over the Str\\\"omgren spheres where bright \\ha\\ emission is present, indicating that the contribution of nebular line emission from the \\hii\\ regions is minimal compared to the scattered stellar continuum and line emission. Therefore, we may use the observed broad-band surface brightness of N30B's reflection nebula and a few simple assumptions to determine the likely position of N30B relative to the B[e] star Hen S22. First, we assume a single value for the spherical albedo of every dust grain.  We also assume that the reflection nebula is optically thick, i.e., every photon from Hen S22 incident upon the reflection nebula is either absorbed or scattered.  It has been shown that dust scattering has an angular dependence and follows the phase function $ \\Phi(\\alpha) = \\frac{\\gamma(1 - g^2)}{4\\pi}(1 + g^2 - 2g \\cos \\alpha)^{-3/2}, $ where $\\alpha$ is the deviation of the photon from the forward direction, $\\gamma$ is the spherical albedo of the dust grains, and $g$ is a measurement of the asymmetry of the phase function as given by the expression, $\\gamma g=\\int\\Phi(\\alpha)\\cos \\alpha~d\\omega$ \\citep{HG41}. We have adopted dust grain properties reported by \\citet{LW76} for wavelengths longward of 3000 \\AA, $\\gamma = 0.7$ and $g = 0.60$. As shown in Figure 6, the scattering angle $\\alpha$ is related to the projected distance $d$ from Hen S22 to the reflection nebula by $\\sin \\alpha = 5~{\\rm pc} / d$. Using the phase function and the brightness of Hen S22, we can calculate the surface brightness of the reflection nebula expected for a given geometry of the reflection.  The broad-band photometry of Hen S22 from 1957 to 1991 show stable brightnesses and colors:  $V = 11.80\\pm0.05$ and $B-V$ = +0.28 \\citep{Zetal86,Zetal96}. The reddening of Hen S22, $E(B-V)$ = 0.25--0.30, is higher than those of the other stars in DEM\\,L\\,106 because of its edge-on disk \\citep{Zetal86}. The broad-band photometry of Hen S22 made by \\citet{oey96a,oey96b} in 1994 shows a similar $V = 11.879\\pm0.004$, but a higher $E(B-V)$ = 0.46. We have adopted an $E(B-V)$ of 0.25 to deredden the Johnson $B$ magnitude of Hen S22, and converted it to the WFPC2 $B$ ($F439W$) magnitude. We then calculate the expected surface brightness of the reflection nebula in the WFPC2 $B$ band for a range of scattering angles. To compare these predictions to observations, a reddening appropriate for the reflection nebula needs to be applied. As the reflection nebula is on the surface of N30B, its extinction should be no greater than the minimum extinction found among stars in N30B, $E(B-V) \\sim 0.06$. We have applied this amount of extinction to the expected surface brightness of the reflection nebula, and the results are presented in Table~3. The maximum expected surface brightness comes at a scattering angle of $\\sim$40$^{\\circ}$, which corresponds to a distance of $\\sim$8~pc between Hen S22 and the reflection nebula. We have measured an average surface brightness of 1.6~$\\times$~10$^{-14}$~ergs~s$^{-1}$~cm$^{-2}$~arcsec$^{-2}$ for N30B's reflection nebula in the  WFPC $B$ band. This observed surface brightness is 20\\% higher than the maximum expected surface brightness. Considering that our assumptions and calculation are crude, the observed and expected surface brightness of the reflection nebula are in remarkably good agreement. As discussed in \\S3.5, the reflection nebula is above the equatorial disk and sees a less reddened Hen S22, the expected surface brightness of the reflection nebula is thus higher and would be in even better agreement with the observations. We conclude that the reflection nebula around N30B is consistent with forward scattering by the dust grains, i.e., N30B is on the near side of Hen S22. The most likely position of N30B is $\\sim$8 pc from Hen S22 and the scattering angle $\\alpha$ is $\\sim$40$^\\circ$.  \\subsection{Periscopic Views of the Polar and Disk Winds of Hen S22}  The broad \\ha\\ line profiles from the reflection nebula around N30B not only change with position (Figure 7) but also differ from that of Hen S22 (Figure 4). Adopting the most likely geometry of reflection suggested in \\S3.4, $d \\sim$ 8 pc and $\\alpha \\sim 40^\\circ$, the light travel time from Hen S22 to us via the reflection nebula is 6 yrs longer than the direct travel time. Thus the spectra of the reflection nebula should be compared to the spectra of Hen S22 taken 6 yrs earlier; unfortunately none are vailable with sufficient resolution.  Spectra of Hen S22 available in the literature were taken from the late 1960's to early 1980's. While these observations yield a consistent systemic velocity of Hen S22, $293\\pm3$ \\kms\\ \\citep{Zetal86}, different blue-shifts (from the systemic velocity) of the absorption component in the hydrogen Balmer lines have been reported: $-$70 \\kms\\ derived from the H8 to H31 lines observed with a reciprocal dispersion of 20 \\AA\\ mm$^{-1}$ in the years 1968--1975 \\citep{M78}, $-$87 \\kms\\ from the \\ha\\ line observed with 38 \\AA\\ mm$^{-1}$ in 1979, and $-$40 \\kms\\ from the H$\\beta$, H$\\gamma$, and H$\\delta$ lines observed with 20 \\AA\\ mm$^{-1}$ in 1982 \\citep{Zetal86}. Since these measurements were made for different Balmer lines with different spectral dispersion at different epochs, it is not clear whether the differences in the blue-shifts of the absorption component are caused by atmospheric effects (low and high Balmer lines have different profiles) or temporal variations in the stellar wind. Our echelle observations of Hen S22 were made in 2000 January and December with a reciprocal dispersion of 3.4 \\AA\\ mm$^{-1}$. These two spectra show almost identical \\ha\\ profiles, with the absorption component detected at 260 \\kms. Relative to a systemic velocity of 293 \\kms, the blue-shift of our \\ha\\ absorption is $-$33 \\kms. To compare our \\ha\\ profiles to that reported by \\citet{Zetal86}, we have convolved our spectra with a Gaussian with $\\sigma$ = 32 \\kms, but we find that the overall \\ha\\ profiles do not match that presented in their Figure 5. This mismatch might be caused by an incorrect wavelength scale in their Figure 5, where the displacement of the absorption dip from the systemic velocity appears to be $<$50 \\kms, significantly less than what they reported in their text.  From the above comparisons we find no conclusive or convincing evidence to indicate that the \\ha\\ line profile of Hen S22 varied. Furthermore, the $V$ magnitude and the $B-V$ color of Hen S22 did not change for more than 0.1 mag from 1957 to 1996 \\citep{Zetal86,Zetal96,oey96b}. Therefore, we cannot attribute the spatial variations of the broad \\ha\\ line from the refection nebula to hypothetical temporal variations of the \\ha\\ line of Hen S22.  Figure 7 shows smooth spatial variations of the broad \\ha\\ profile, especially the absorption component, across N30B, similar to what has been observed in Car II \\citep{Betal98}. The explanation of this spatial variation is best illustrated by the Homunculus nebula of $\\eta$ Car, as described and analyzed in detail by \\citet{smith02}. The bipolar lobes of the Homunculus are dusty and scatter the incident starlight of $\\eta$ Car; the reflected spectra at higher latitudes show \\ha\\ absorption extending to larger blue-shifts. As the Homunculus is smaller than 1 ly in size, temporal variations in the spectrum of $\\eta$ Car are negligible over the differential light travel time; therefore, the latitude-dependence of the \\ha\\ absorption in the reflected spectra must indicate an anisotropic stellar wind from $\\eta$ Car. Hen S22 has been shown to possess both a slow disk wind and a fast polar wind \\citep{Zetal86}, consequently its spectral properties should also be latitude-dependent. Thus, the variations in the absorption component in the reflected \\ha\\ emission at N30B must be a result of viewing Hen S22 from different latitudes, too. The reflected spectra on the west side of N30B resemble better the observed \\ha\\ line of Hen S22, while those on the east side show the absorption component extending to larger blue-shifts, almost $-$100 \\kms\\ from Hen S22's systemic velocity of 293 \\kms. These variations suggest that the east side of N30B views the wind of Hen S22 from a higher latitude and the west side of N30B is close to the disk plane.  The absorption component in the \\ha\\ line of a B[e] star arises from the dense, disk component of the wind close to the star, as illustrated in the cartoon in Figure 7 of \\citet{Zetal85}. The variations of the absorption component in the reflected stellar \\ha\\ line emission across N30B thus offer an invaluable opportunity to study the transition from the disk component to the polar component of the stellar wind of a B[e] star. Future high-S/N, high-resolution spectroscopic observations and stellar atmospheric modeling are needed for a better understanding of the stellar wind of the hybrid-spectrum B[e] supergiant Hen S22."
    },
    "0212/astro-ph0212524_arXiv.txt": {
        "abstract": "Light curves  of eight BL Lac objects in the BVRI bands have been analyzed. All of the objects tend to be bluer when brighter. However spectral slope changes differ quantitatively from those of a sample of QSOs analyzed in a previous paper \\citep{tre02} and appear consistent with a different nature of the optical continuum. A simple model representing the variability of a synchrotron component can explain the spectral changes. Constraints on a possible thermal accretion disk component contributing to the optical luminosity are discussed. ",
        "introduction": "Multi-wavelength observations of active galactic nuclei (AGN) suggest the presence of various emission components with different relative weights in different AGN classes. The large number of parameters involved in physical models of  emission mechanisms can hardly be constrained by single epoch observations. Multi-band variability studies are then crucial in identifying and constraining individual components. Moreover, active nuclei are not spatially resolved in general. For this reason variability time scales, delays and correlations between variations in different continuum components or emission lines provide the main  information about the location of different components. The very definition of  AGN classes is in part based on variability properties. In fact, while quasars (QSOs) and Seyfert galaxies vary, on average,  less than 0.5 mag on time scales of months in the optical band, blazars (i.e. BL Lac objects and Optically Violent Variables, OVVs) may change their luminosities by more than one mag in a few days \\citep{ulr97}. Quantifying spectral variability will then help understanding different AGN classes by comparison. Although monitoring campaigns, from radio to X-ray wavelengths, exist for a limited number of objects, they are difficult to realize with adequate time sampling and duration. For this reason we try to exploit multi-band variability studies limited to  the optical range, which in some case are available for several objects and relatively long observing periods. In a previous paper \\citet{tre02} analyzed the spectral slope variability of a sample of PG QSOs on the basis of the light curves in the B and R bands, provided by the Wise Observatory group, who monitored 42 quasars for 7 years with a typical sampling interval of about one month \\citep{giv99}. A spectral variability parameter $\\beta \\equiv \\Delta \\alpha / \\Delta \\log f_{\\nu}$, $f_{\\nu}$ being  the specific flux and $\\alpha\\equiv \\partial \\log f_{\\nu} / \\partial \\log \\nu$ the spectral slope, was defined by \\citet{tre02} who derived constraints on variability mechanisms from the average $\\alpha$ and $\\beta$ values. In particular they showed that the variation of spectral slope implied by changes of the accretion rate are not able to explain $\\beta$ values deduced from the observations, while ``hot spots'' on the accretion disc, possibly caused by instability phenomena \\citep{kaw98}, can easily account for the observed spectral variability. In the present study we analyze B, V, R, I light curves of eight BL Lac objects obtained by \\citet{dam02} in a monitoring campaign of about 5 years with an average sampling interval of  $\\approx 25$  days. We show that  BL Lacs clearly differ from QSOs in their $\\alpha,\\beta$ distribution. We  show that a simple model representing the variability of a synchrotron component can account for the observed $\\alpha$ and $\\beta$ values and we discuss  constraints which can be derived from spectral variability on the amount of thermal radiation contributed by a disk component in the optical band. ",
        "conclusions": "We have performed a new analysis of the light curves of eight BL Lac objects monitored in B,V,R,I bands for a total period of about 5  years with an average sampling interval of $\\approx 25$ days.  Adopting a spectral variability index $\\beta$, representing the average ratio between the change of spectral slope and the logarithmic luminosity change, we find that BL Lacs show smaller spectral variability with respect to quasars (analyzed in a previous paper). This happens despite the fact that both the spectral slope and luminosity variations of BL Lacs are larger then those of QSOs. BL Lacs and QSOs appear clearly segregated in the $\\alpha-\\beta$ plane. Our analysis allows a quantitative comparison of the observations with a model based on variability of the synchrotron component. The model easily accounts for the observed position in the $\\alpha-\\beta$ plane with a natural choice of the relevant parameters. Thus the segregation in the $\\alpha-\\beta$ plane is a consequence of the different emission mechanism in the optical band: synchrotron in the case of BL Lacs and thermal hot spots on the accretion disk in the case of QSOs \\citep[see][]{tre02}. A strong thermal bump can produce negative $\\beta$ values, which are not observed in our sample, while they do in fact appear in some typical LBL object, where the thermal bump emerges in the faint phases. Thus spectral variability, even restricted to the optical band, can be used to set limits on the relative contribution of the synchrotron component and the thermal component to the overall SED. Further studies of the spectral variability index $\\beta$ in different objects will allow us to establish whether the thermal bump is significant only in the most extreme LBL, or to what extent its contribution, possibly correlated with the accretion rate $\\dot m$, is related to total luminosity."
    },
    "0212/astro-ph0212238_arXiv.txt": {
        "abstract": "We present preliminary results from a study of the neutral hydrogen (\\hi{}) properties of an X-ray selected sample of nearby loose galaxy groups.  This forms part of a multi-wavelength investigation (X-ray, optical and radio) of the formation and evolution of galaxies within a group environment.  Some initial findings of an ATNF Parkes Multibeam wide-area neutral hydrogen (\\hi{}) imaging survey of 17 nearby galaxy groups include two new, potentially isolated clouds of \\hi{} in the NGC~1052 and NGC~5044 groups and significant amounts of \\hi{} within the group virial radii of groups NGC~3557 and IC~1459 - two groups with complex X-ray structures that suggest they may still be in the act of virialisation.  Here we present ATCA high-resolution synthesis-imaging follow-up observations of the distribution and kinematics of \\hi{} in these four groups. ",
        "introduction": "Neutral hydrogen is a valuable tracer of galactic structure and dynamics. It is the most spatially extended component of a galaxy's disk, and is particularly sensitive to tidal interactions and mergers between galaxies, retaining the imprint of an encounter out to large galactic radii \\cite{mpa+95}. It is therefore an ideal tool for probing the morphology and dynamics of galaxies in groups, with the aim of unravelling the processes driving galaxy evolution in a group environment.  This study is part of the GEMS project - a multi-wavelength (X-ray, optical and radio) investigation of the formation and evolution of galaxies within groups We have conducted a wide-area ($\\degr{5.5} \\times \\degr{5.5}$ per target) neutral hydrogen~(\\hi{}) imaging survey of 17 nearby galaxy groups (distance $< 40$~Mpc) using the ATNF Parkes Radiotelescope Multibeam System, at an angular resolution of 17~arcminutes.  Initial results of this survey include the detection of \\hi{} within the X-ray virial radius in the IC 1459 and NGC~3557 groups. Additionally, new group members previously uncatalogued at optical wavelengths were detected at the edges of the groups NGC~1052 and NGC~5044.  High angular resolution ($\\sim 35$~arcsec) synthesis followup observations of a targeted 33~arcmin region in each of these four groups were subsequently obtained, using the Australia Telescope Compact Array, with the goal of determining the precise location, distribution and kinematics of the \\hi{} in these regions. ",
        "conclusions": "In combination with forthcoming results from the remaining 13 groups in this study, and in parallel with the work being carried out on this sample at X-ray and optical wavelengths, these results will help us to understand how the neutral hydrogen gas in loose galaxy groups, through its interactions with other group elements, contributes to the evolution of the galaxies and the groups they inhabit."
    },
    "0212/astro-ph0212476_arXiv.txt": {
        "abstract": "In view of the current efforts to extend the KASCADE experiment (KASCADE-Grande) for observations of Extensive Air Showers (EAS) of primary energies up to 1$\\cdot 10^{18}$ eV, the features of muon arrival time distributions and their correlations with other observable EAS quantities have been scrutinised on basis of high-energy EAS, simulated with the Monte Carlo code CORSIKA and using in general the QGSJET model as generator. Methodically various correlations of adequately defined arrival time  parameters with other EAS parameters have been investigated by invoking non-parametric methods for the analysis of multivariate distributions, studying the classification and misclassification probabilities of various observable sets. It turns out that adding the arrival time information and the multiplicity of muons spanning the observed time distributions has distinct effects improving the mass discrimination. A further outcome of the studies is the feature that for the considered ranges of primary energies and of distances from the shower axis the discrimination power of global arrival time distributions referring to the arrival time of the shower core is only marginally enhanced as compared to local distributions referring to the arrival of the locally first muon. ",
        "introduction": "Since the first experimental studies in 1953 by Bassi, Clark and Rossi \\cite{01} and Jelley and Whitehouse \\cite{02} arrival time distributions of the charge particle components of Extensive Air Showers (EAS) have been often theoretically and experimentally investigated under various aspects \\cite{03,04,05,06,07,08,09}. Systematic studies are performed particularly at large EAS detector installations: in Haverah Park \\cite{10,11,12,13,14,15,16}, at the Potchefstroom University \\cite{17}, in Akeno and on Mt. Chacaltaya \\cite{18,19,20}, at MSU \\cite{21}, more recently with the GREX/ COVER-PLASTEX \\cite{22,23,24,25,26} and KASCADE \\cite{27,28,29,30,31,32,33,34} experiments. Measurements of arrival time distributions of high energy hadrons near the shower core have been reported by the Maryland University group \\cite{35} and by KASCADE \\cite{36}. In addition to the fact that the observed arrival time distributions of the EAS particles display the phenomenological appearance (shape and structure) of the EAS disc, they provide also a coded picture of the longitudinal EAS development. Especially the muon component, when multiple scattering in the atmosphere at higher muon energies gets negligible, maps the distributions of the production heights via the time-of-flight of the muons from the loci of decay of the charged parent pions to the detector \\cite{27,37}. Fig.~1 sketches the relation in a simplified way. Muons released in higher atmospheric altitudes (and observed on ground at  fixed larger distances from the shower axis) show smaller delays relative to the arrival of the shower centre (\"light front\"), and their (relative) arrival times let expect some discrimination features for the mass of the EAS primary \\cite{05,06,07,27,29}. Thus the arrival time distributions of muons, produced in iron induced showers e.g. should get shifted to shorter delays due to the faster development as compared to proton induced showers of the same primary energy. Of course, a serious analysis of this mapping has to invoke realistic simulations of the EAS development and of the muon tracking, taking into account the influence of multiple scattering and off-axis production \\cite{37} of the muons. Fig. 1 indicates also that observations of the angle distributions of the muon incidence provide alternative experimental possibilities \\cite{38,39}, in case that multiple scattering effects do not seriously obscure that information \\cite{37}. The information content of the combination of  both types of EAS observations, known as \"Time-Track-Complimentary-Principle\" \\cite{37,40}, has been scrutinised in Ref. \\cite{29}.  \\begin{figure}[!t] \\centering \\includegraphics[bb=187 202 492 675, angle=270, width=15.0cm]{fig1.ps} \\caption{\\label{Fig. 1} Sketch of a simplified relation between muon arrival time and angle of muon incidence, respectively, with the production height, neglecting multiple scattering and off-axis production ($\\beta _{\\mu} c$ = velocity of the muon). Additionally (right) the definitions of the arrival time quantities characterising the shape and the structure of the muon disc are indicated.} \\end{figure}  Arrival times $\\tau ^{1}_\\mu < \\tau ^{2}_\\mu < \\tau ^{3}_\\mu$,... of the EAS muons, locally registered by timing detectors at a distance $R _\\mu$ from the shower axis have to refer to a well defined zero-time, usually the arrival time $\\tau _{c}$ of the shower core (global arrival times):  \\begin{equation*} \\Delta \\tau ^{n\\ glob} _\\mu = \\tau ^n _\\mu (R_\\mu) - \\tau _c \\end{equation*}  e.g. \\begin{equation*} \\Delta \\tau ^{1\\ glob}_\\mu = \\tau ^1 _\\mu (R_\\mu) - \\tau _c \\end{equation*}  Often there are experimental difficulties  to determine the arrival time $\\tau _c$ with sufficient experimental precision, and only \"local\" times are considered. They refer to the foremost (first locally registered) muon:  \\begin{equation*} \\Delta \\tau ^{n\\ loc} _\\mu (R_\\mu) = \\tau ^n _\\mu (R_\\mu) - \\tau ^1 _\\mu (R_\\mu) \\end{equation*}  {\\it Local} arrival time distributions display the internal structure of the shower disc (characterised by various time parameters), but they do not carry information about the shape (curvature) of the front. Nevertheless studies based on simulations of the EAS development \\cite{16,19} have shown that mass discrimination effects are just pronounced in observations of $ \\Delta \\tau ^{1\\ glob} _\\mu $.  For event-by-event observations with a fluctuating number of muons, the single relative arrival time distributions can be characterised by the mean values $\\Delta \\tau _{mean}$, and by various quantiles $\\Delta \\tau _q$, like the median $\\Delta \\tau _{0.50}$, the first quartile $\\Delta \\tau _{0.25}$ and the third quartile $\\Delta \\tau _{0.75}$. For ordered statistics of measured times $\\Delta \\tau _1 \\le \\Delta \\tau _2 \\le . . . \\le \\Delta \\tau _n$ and $k := n \\alpha + \\xi, k$ integer and $\\xi \\in [0,1)$, the $\\alpha$-quantile $\\Delta \\tau _\\alpha$ ~ $(\\alpha \\in (0,1))$ is the following \\cite{28,29}: \\begin{equation*} \\Delta \\tau _\\alpha = \\left\\{ \\begin{array}{ll} ( \\Delta \\tau _k + \\Delta \\tau _{k+1} ) / 2 : & for \\: \\: \\xi = 0 \\\\ \\Delta \\tau _k : & for \\: \\: \\xi \\in (0,1) \\end{array} \\right. \\end{equation*} That means: in the case of large $n$, a fraction $\\alpha$ of muons has arrival times less than $\\Delta \\tau _\\alpha$. The {\\it mean values} and {\\it dispersions} (standard deviations) represent the time profile of the EAS muon component. Measurements of muon arrival time distributions and the determination of the distributions of various different time quantities have been a subject of the current investigations of the KASCADE experiment \\cite{27,28,30,31,32,33,34}. In these investigations the temporal EAS structure has been studied in detail in dependence on the shower size $N_e$, the muon number $N_\\mu$, or the truncated muon number $N _\\mu ^{tr}$ (which is used as approximate energy identifier in KASCADE observations \\cite{41}), respectively, and from the angle of EAS incidence $\\theta$. Special features arising from the observed muon multiplicity have been revealed \\cite{33}. More recently \\cite{34} using larger samples of shower events, experimentally observed with KASCADE,  the sensitivity of local muon arrival time distributions and of their correlations with other EAS observable to the mass of the EAS primaries has been investigated. Methodically advanced statistical analysing techniques have been applied \\cite{42}, based on Bayesian decision making for the classification of the EAS events. The procedure requires simulated distributions (processed through the detector response) as reference patterns, provided by the Monte Carlo simulation program CORSIKA \\cite{43} with a particular model of the hadronic interactions. Thus the analysis necessarily implies some model dependence. But the consistency with the invoked QGSJET model \\cite{44} could be established thanks to the measurements with varying distance $R_\\mu$ from the shower centre and varying muon multiplicity thresholds for being accepted in the observation sample. Nevertheless, it has been also shown that the mass discrimination power of local muon arrival time distributions is rather marginal in the observed range $R_\\mu < 100$ m and primary energies $E_0 \\le 10^{16}$ eV, and it does not significantly help for the classification.  The phenomenological features of the time structure of high energy EAS at larger distances from the shower axis have been experimentally studied with the Haverah Park \\cite{10,11,12,13,14,15,16} and the Akeno air shower arrays \\cite{18,19,20}, and the results have been compared with phenomenological model predictions (scaling models and multiplicity prescriptions for the particle production in hadronic collisions). An analysis of the mass discrimination power on basis of modern models of the hadronic interactions has not yet been performed for these cases.  The recently started extension of the KASCADE detector installation: KASCADE-Grande \\cite{45} will not only provide the possibility to extend the studies of muon arrival time distributions and their correlations to larger distances $R_\\mu$ and larger energies of the observed EAS. There is also a chance to measure global arrival times, at least for $R_\\mu$ up to 350 m \\cite{46}. The investigations of this paper use the analysis techniques of our previous studies \\cite{29,34} and refer mainly to the QGSJET \\cite{44}, in some few comparisons also to other models. The studies are focused to explore the basic information content of muon arrival time distributions of EAS with energies  up to ca. 1$\\cdot 10^{18}$ eV. The analyses are completely based on data of simulated showers and consider ideal cases, i.e. neglecting the influence of the detection system. The view of interest is the discrimination of the mass of the primary particles (with adopting a particular hadronic model as generator of the reference patterns). For some observable combinations the differences in the predictions of Monte Carlo simulations using different high-energy models as generators have been scrutinised. Only marginal differences for the arrival time distributions have been found. ",
        "conclusions": "In the present paper based on realistic EAS Monte Carlo simulations the role of muon arrival time distributions has been scrutinised in view of correlations with the main EAS observables to be measured with the KASCADE-Grande layout for the discrimination of the primary mass at higher primary energies. The main observables for that purpose are the total number of the charged particles $N_{ch}$ and a part of the EAS muon intensity $N_\\mu ^{part}(R_\\mu)$ registered with the KASCADE array, embedded in KASCADE-Grande. The partial muon number $N_\\mu ^{part}(R_\\mu)$ is dependent from  the location of the EAS centre. For the present simulation studies it has been replaced by the density $\\rho _\\mu (R_\\mu)$ ($E_\\mu > 240$ MeV). In summary, the analyses of the classification and misclassification probabilities, determined with non-parametric statistical techniques give evidence that the information from muon arrival time distributions has a potential for improving the mass discrimination. This conclusion, being in contrast to findings at lower energies, holds for global time quantities as well as for the local ones. Actually the EAS profiles of the two different types of time parameters differ by an offset being only marginally dependent from the mass. This is due to an increasing ``flatness'' of the $<\\Delta \\tau _1>$ profile with increasing primary energy, and differences are moved to larger distances out of experimental reach. This feature is important with respect to the experimental difficulties to define $\\tau _c$ and to measure global time quantities. Obviously in the considered $R_\\mu$ range the local times provide most of the accessible information.  The measuring techniques of muon arrival time distributions with the Central Detector of KASCADE imply the simultaneous determination of the multiplicity $n$ of muons spanning the arrival distributions. This multiplicity is related to the density $\\rho_\\mu ^* (R _\\mu)$ of the higher energy muons detected with the facilities of the Central Detector. The quantity $\\rho _\\mu ^* (R _\\mu )$ (or the  number of muons experimentally detected with the timing facility), when correlated with timing measurements (or used as $\\Delta \\tau _q ^* = \\Delta \\tau _q (R _\\mu /\\rho _\\mu ^* (R_\\mu)$) plays an important role in improving the effects of the arrival time observations.  Additionally (as outlined in Ref.\\cite{34}) as consequence of the specific observation conditions muon arrival time observations establish a selected subset of all showers, and the determination of the mass composition needs a corresponding correction \\cite{60}. Thus the results inferred with a variation of the multiplicity threshold $n_{thr}$ and of the observation distance $R_\\mu$ provide a consistency test for the model used for the non-parametric analyses of the observed EAS.  It should be noted that the present results refer to ideal cases since any detector response functions, detector efficiencies and  limitations of the statistical accuracies due to limited detector sizes  have been ignored. Additionally the necessarily limited number of simulated EAS implies a limit on the reasonable number of observables used in a single analyses.  \\ack  The Deutsche Forschungsgemeinschaft and the Centre of Excellence IDRANAP in the National Institute for Physics and Nuclear Engineering, Bucharest have considerably supported the present studies. Some of us (I. M. B., H. R. and C. A.) would like to thank Prof. Dr. H. Bl{\\\"u}mer and Prof. Dr. D. Poenaru for the kind hospitality in Forschungszentrum Karlsruhe and in National Institute for Physics and Nuclear Engineering - Horia-Hulubei, Bucharest, respectively, experienced during various mutual research visits related to the reported studies."
    },
    "0212/astro-ph0212195_arXiv.txt": {
        "abstract": "The effect of the shock propagation on neutrino oscillation in supernova is studied paying attention to evolution of average energy of $\\nu_{e}$ and $\\bar{\\nu}_{e}$. We show that the effect appears as a decrease in average $\\nu_{e}$ (in case of inverted mass hierarchy, $\\bar{\\nu}_{e}$) energy at stellar surface as the shock propagates. It is found that the effect is significant 2 seconds after bounce if $3 \\times 10^{-5} < \\sin^{2}{\\theta_{13}} < 10^{-2}$. ",
        "introduction": "Recently effects of shock propagation on neutrino oscillation in supernova was studied \\cite{SchiratoFuller2002,Lunardini2003} and it was shown that some characteristic signatures emerge as the shock propagates through the regions where matter-enhanced neutrino flavor conversion occurs.  There have been many studies on neutrino oscillation in supernova: extracting information of neutrino parameters from the observation of SN1987A neutrinos \\cite{Smirnov1994,Jegerlehner1996,MinakataNunokawa2001,Lunardini2001} or a future supernova neutrinos \\cite{Dighe2000,KT1,KT4,KT2,KT3}, and probing supernova physics from observed neutrinos of a future supernova \\cite{Minakata2001}. But all of them are done without the effect of the shock propagation.  In this paper the effect of the shock propagation is studied paying attention to evolution of average energies of $\\nu_{e}$ and $\\bar{\\nu}_{e}$. We show when and with which parameter ($\\sin^{2}{2\\theta_{13}}$) the effect is significant or can be neglected safely. ",
        "conclusions": "As we saw in the previous section, neutrino average energy decrease in general as the shock propagates. This is because the shock propagation cause decrease in the adiabaticity parameter of H-resonance and suppress the conversion between flavors. But it should be noted that the original spectra will change due to the evolution of the protoneutron star and the neutrinosphere. Thus we can not say about the value of $\\sin^{2}{2\\theta_{13}}$ and the mass hierarchy only from the evolution of the average energy of the observed neutrino. To do so, we need a model of supernova and spectrum evolution.  In this paper we studied the effect of the shock propagation paying attention to evolution of average energies of $\\nu_{e}$ and $\\bar{\\nu}_{e}$. It is shown that the effect appears as a decrease in average $\\nu_{e}$ (in case of inverted mass hierarchy, $\\bar{\\nu}_{e}$) energy at stellar surface as the shock propagates. Further it is found that the effect is significant 2 seconds after bounce if $3 \\times 10^{-5} < \\sin^{2}{\\theta_{13}} < 10^{-2}$."
    },
    "0212/hep-ph0212005_arXiv.txt": {
        "abstract": "We analyze the curvaton scenario in the context of supersymmety. Supersymmetric theories contain many scalars, and therefore many curvaton candidates. To obtain a scale invariant perturbation spectrum, the curvaton mass should be small during inflation $m \\ll H$. This can be achieved by invoking symmetries, which suppress the soft masses and non-renormalizable terms in the potential. Other constraints on the curvaton model come from nucleosynthesis, gravitino overproduction, and thermal evaporation.  The curvaton coupling to matter should be very small to satisfy these constraints ({\\it e.g.} $h \\lesssim 10^{-8}$ for typical soft masses $m \\sim \\TeV$). ",
        "introduction": "It is now widely believed that the early universe went through a period of superluminal expansion, known as inflation.  In addition to explaining the isotropy and homogeneity of our observable universe, inflation can generate the seeds for structure formation.  In the usual picture the quantum fluctuations of the slowly rolling inflaton field ``freeze in'' soon after horizon exit, and become essentially a classical perturbation which remains constant until the moment of horizon re-entry.  The produced perturbations only depend on the potential during inflation, and therefore severely constrain possible inflationary models.  Recently there has been a revived interest in the idea that not the perturbations in the inflaton field are responsible for the cosmic microwave background (CMB) fluctuations, but instead isocurvature fluctuations in another scalar field, the ``curvaton'' field, that is sub-dominant during inflation. After inflation the isocurvature fluctuations of the curvaton field are converted into adiabatic ones. If at decay the contribution of the curvaton to the energy density is considerable, this can lead to perturbations of the observed size. The generation of curvature does not depend on the nature of inflation, beyond the requirement that the Hubble constant must be nearly constant.  This mechanism of generating perturbations was first described years ago~\\cite{early}, but it did not attract much interest until recently~\\cite{sloth,wands,moroi,curvaton,decouple}.  The usual set up is the following~\\footnote[2] {For an alternative implementation of the curvaton model see~\\cite{mar}} : During inflation the curvaton field acquires a large expectation value if it has a negative mass squared, or if its mass is much smaller than the Hubble constant. In the post-inflationary epoch the Hubble expansion acts as a friction term in the equations of motion, and the field remains effectively frozen at a large field value. This stage ends when the Hubble constant becomes of the order of the curvaton mass, $H \\sim m_\\phi$, at which point the curvaton starts oscillating in the quadratic potential.  The curvaton behaves as cold matter; its energy density red shifts as $\\rho_\\phi \\propto a^{-3}$, with $a$ the scale factor of the universe. After inflaton decay, the universe becomes radiation dominated, with the energy density in radiation red shifting as $\\rho_I \\propto a^{-4}$. During this state of mixed matter, {\\it i.e.}, radiation and cold dark matter, the isocurvature perturbations are transformed into adiabatic ones.  The perturbations grow until the curvaton becomes to dominate the universe, or if this never happens, until the curvaton decays.  In this paper we reanalyze the curvaton scenario in the context of supersymmetry (SUSY). SUSY theories contain many scalar fields, and therefore many possible curvaton candidates.  However, the finite energy density in the early universe breaks supersymmetry, and induces soft masses of the order of the Hubble constant for all scalar fields~\\cite{softmass,flat}. Such heavy fields lead to a perturbation spectrum with a large spectral tilt, in contradiction with observations.  Symmetries can be invoked to protect scalars from large soft masses.  In D-term inflation Hubble induced soft masses are forbidden by gauge invariance~\\cite{dterm}. In no-scale gravity models, and generalizations thereof, there is a Heisenberg symmetry that protects scalar fields from soft mass terms at tree level~\\cite{noscale}.  Another possibility is that the curvaton field is a (pseudo) Goldstone boson.  Non-renormalizable terms in the potential can lift the flat potential for the large field values needed in the curvaton scenario.  For the curvaton model to work such terms should be very small or absent up to some high order. Here SUSY helps, since the non-renormalization theorem assures that operators absent in the superpotential are absent at all scales.  There are several constraints on the curvaton model. The produced perturbations should be adiabatic (which requires that the curvaton decays after inflaton decay), nearly scale invariant (the curvaton mass should be much smaller than the Hubble constant during inflation), and of the correct magnitude.  For the perturbations to be sizable, the curvaton energy density needs to be a significant fraction of the total energy density in the universe.  To avoid gravitino overproduction the reheat temperature should be sufficiently low.  Curvaton decay should not destroy the nucleosynthesis predictions. Furthermore, thermal damping and evaporation should be taken into account.  In the next section we discuss these various constraints on the curvaton model in more detail. In section \\ref{mass} we consider the possible (soft) mass terms that can arise in SUSY theories, and analyze the parameter space for the various cases.  We end with conclusions. ",
        "conclusions": "The finite energy density during inflation breaks supersymmetry, and induces soft masses of the order of the Hubble constant for all scalars.  The inflationary perturbation spectrum for these fields is highly non-Gaussian, and therefore they cannot play the role of the curvaton.  To avoid this conclusion one has to invoke symmetries, to keep the soft mass terms small during inflation.  In no-scale type gravities, scalar masses are induced radiatively and are suppressed to the soft breaking scale ($M_s \\sim H$) by loop factors. These theories have the additional advantage that the gravitino mass can be arbitrarily large, thereby avoiding problems with gravitino overproduction. In $D$-term inflation gauge symmetry forbids soft Hubble-induced masses. At the end of inflation, when $D$-terms are converted to $F$ terms and kinetic energy, a soft mass term $|\\delta m^2| \\sim H^2$ appears.  Pseudo-Goldstone bosons are protected from soft mass terms. Their mass is set by the amount of symmetry breaking, and can be small.  Apart from constraints considering gravitino production, the analysis for this case also applies to standard model fields.  In all cases the parameter space for the curvaton scenario is constrained by bounds from nucleosynthesis and early thermal evaporation.  To avoid thermal evaporation the coupling to matter should be small. In particular, for typical low energy soft masses $m \\sim \\TeV$ and an interaction term in the scalar potential of the form $V \\ni h \\phi^2 \\chi^2$, thermal evaporation is negligible for $h \\lesssim 10^{-8}$. But we note that thermal constraints can be circumvented if the inflaton and curvaton sector decouple completely~\\cite{decouple}.   Strong constraints also come from the existence of non-renormalizable operators in the potential, which are suppressed by some cutoff scale $M$.  For operators of the form $ V \\ni \\phi^{n+4}/\\mpl^n$, the curvaton mass is constraint to $m \\lesssim 10^3 \\GeV$ for $n=2$, $m \\lesssim 10^7 \\GeV$ for $n=4$, and $m \\lesssim 10^8 \\GeV$ for $n=6$. These bounds are absent for $D$-term inflation, with a negative Hubble induced mass squared after inflation.  One can ask whether there are any natural candidates for the curvaton field.  Affleck-Dine condensates have masses of the low energy SUSY breaking scale, $m \\sim \\TeV$. Along most flat directions the $\\phi^6$ term is absent. However, couplings to matter are too large, and do not satisfy the curvaton constraint $h < 10^{-8}$.  Moduli and polyoni fields have Planck mass suppressed couplings and typically decay after nucleosynthesis.  This can be avoided if the mass term is sufficiently high, $m \\gtrsim 100 \\TeV$.  In no-scale gravities this is naturally achieved by giving them a mass of the order of the gravitino mass.  But that means that also their mass during inflation is too large ($\\delta m^2 \\sim H^2$), and they cannot be the curvaton.  Right-handed sneutrinos can have small couplings to matter. In the sea saw model, neutrinos obtain masses $m_\\nu \\sim h \\langle H_u \\rangle/M_N \\lesssim 10^{-3} \\eV$, where $\\langle H_u \\rangle$ is the expectation value of the Higgs field and $M_N$ the sneutrino mass.  In \\cite{lepto} leptogenesis from an sneutrino dominated universe is considered.  In this model, sneutrino has mass $M_N \\sim 10^7 \\GeV$ and $h \\lesssim 10^{-12}$. The sneutrino can be the curvaton if there are no constraints from gravitino production.  Moreover, if the sneutrino mass is generated at some GUT scale, one needs to explain why $M_N \\ll M_{GUT}$ and why non-renormalizable operators are absent to very high order.  This work was supported by the European Union under the RTN contract HPRN-CT-2000-00152 Supersymmetry in the Early Universe."
    },
    "0212/astro-ph0212310_arXiv.txt": {
        "abstract": "Based on observed rotation curves of galaxies and theoretical simulations of dark matter halos, there are reasons for believing that at least three different types of dark matter halos exist in the Universe classified by their masses $M$ and the inner slope of mass density $-\\alpha\\,$: Population A (galaxies): $10^{10} h^{-1} M_\\odot \\la M \\la 2\\times 10^{13} h^{-1} M_\\odot$, $\\alpha \\approx 2$; Population B (cluster halos): $M \\ga 2\\times 10^{13} h^{-1} M_\\odot$, $\\alpha \\approx 1.3$; and Population C (dwarf halos): $M \\la 10^{10} h^{-1} M_\\odot$, $\\alpha \\approx 1.3$. In this paper we calculate the lensing probability produced by such a compound population of dark halos, for both image separation and time delay, assuming that the mass function of halos is given by the Press-Schechter function and the Universe is described by an LCDM, OCDM, or SCDM model. The LCDM model is normalized to the {\\em WMAP} observations, OCDM and SCDM models are normalized to the abundance of rich clusters. We compare the predictions of the different cosmological models with observational data and show that, both LCDM and OCDM models are marginally consistent with the current available data, but the SCDM model is ruled out. The fit of the compound model to the observed correlation between splitting angle and time delay is excellent but the fit to the number vs splitting angle relation is only adequate using the small number of sources in the objective JVAS/CLASS survey. A larger survey of the same type would have great power in discriminating among cosmological models. Furthermore, population C in an LCDM model has a unique signature in the time domain, an additional peak at $\\sim 3$ seconds potentially observable in GRBs, which makes it distinguishable from variants of CDM scenarios, such as warm dark matter, repulsive dark matter, or collisional dark matter. For image separations greater than 10 arcseconds the differently normalized LCDM models predict significantly different lensing probabilities affording an additional lever to break the degeneracies in the CMB determination of cosmological parameters. ",
        "introduction": "It is well known that gravitational lensing is a powerful tool for directly probing the structure and distribution of dark matter in the Universe \\citep[and references therein]{tur84,sch92,bar01,cou02}. By comparing the number of lenses found in a survey of remote sources (e.g., quasars, radio galaxies, or high redshift type Ia supernova) to theoretical predictions, we should be able to deduce the quantity of dark matter in the Universe and how it is distributed \\citep[henceforth LO02; Gladders et al. 2003]{tur90,wam95,por00,kee01a,kee01,li02}. The joint observations of gravitational lensing, high redshift type Ia supernova, cosmic microwave background (CMB), and cluster abundances constrain the Universe to be in all likelihood flat and accelerating, with the present mass density being composed of about 70\\% cosmological constant (or dark energy), 26\\% dark matter, and 4\\% ordinary matter \\citep{ost95,bah99,wan00,pea01,efs02,yam02,spe03}.  However, the lensing cross-section (and thus the lensing probability) is found to be extremely sensitive to the inner density profile of lenses (Keeton \\& Madau 2001; Wyithe, Turner, \\& Spergel 2001; LO02). For example, with fixed total mass, when the inner slope of the density profile, $- \\alpha$, changes from $-1$ [the NFW case \\citep{nav96,nav97}] to $-2$ [the singular isothermal sphere (SIS) case \\citep{got74,tur84}] while maintaining the same mass density in lenses, the integral lensing probability increases by more than two orders of magnitudes for the flat model of the Universe (LO02). Therefore, lensing also sensitively probes small scale structure. This complicates matters and renders it is hazardous to use observed lensing statistics to draw inferences with regard to cosmology before determining the sensitivity to other factors.  In LO02, we have shown that in order to explain the observed numbers of lenses found in the JVAS/CLASS survey, at least two populations of dark halos must exist in nature. One population, which corresponds to normal galaxies, has masses $\\la 10^{13} M_\\odot$ and a steep inner density profile ($\\alpha \\approx 2$, i.e. SIS) presumably determined by the distribution of baryonic material in the inner parts of galaxies;\\footnote{Strictly speaking, the mass density profile of galaxies is the combination of the density profile of hark halos and that of baryonic material at the centers. For brevity, in the paper we still call it the mass density of halos, though we mean the sum of the mass density of dark halos and that of baryonic material whenever we refer to galaxies.}  the other one, which corresponds to groups or clusters of galaxies, has masses $\\ga 10^{13} M_\\odot$ and a shallow inner density profile ($\\alpha \\la 1.4$, i.e. similar to NFW). A similar conclusion has been obtained by \\citet{por00} for explaining the number of lenses found in the CASTLES survey. These results are consistent with the theoretical studies on the cooling of massive gas clouds: there is a critical mass of halos $\\sim 10^{13} M_\\odot$ below which cooling of the corresponding baryonic component will lead to concentration of the baryons to the inner parts of the mass profile \\citep{ree77,blu86}.  In this paper we investigate the lensing statistics produced by a compound population of halos. We assume that there are three populations of halos in the Universe:  \\noindent --- Population A: $10^{10} h^{-1} M_\\odot < M < 2\\times 10^{13} h^{-1} M_\\odot$, $\\alpha = 2$ (SIS); \\\\ --- Population B: $M > 2\\times 10^{13} h^{-1} M_\\odot$, $\\alpha = 1.3$ [GNFW \\citep[generalized NFW,][]{zha96}]; \\\\ --- Population C: $M < 10^{10} h^{-1} M_\\odot$, $\\alpha = 1.3$ (GNFW),  \\noindent where $h$ is the Hubble constant in units of 100 km s$^{-1}$ Mpc$^{-1}$. Population A corresponds to spiral and elliptical galaxies, whose centers are dominated by baryonic matter. Population B corresponds to groups or clusters of galaxies, whose centers are dominated by dark matter. Population C corresponds to dwarf galaxies or subgalactic objects, whose centers lack baryons due to feedback processes such as supernova explosions, stellar winds, and photoionizations \\citep{kau93,som99,ben02,ma03}, and so are also dominated by dark matter. We adopt an inner slope for the dark matter halos of $\\alpha = 1.3$ consistent with the value $1.3\\pm 0.2$ found by \\citet{sub00} and intermediate between the values advocated by \\citet{nav96,nav97} of $\\alpha = 1.0$ and \\citet{ghi00} of $\\alpha = 1.5$. We will calculate here the lensing probability of two measurable variables: image separation and time delay, examining in a subsequent paper the expected arc properties.  Recently, \\citet{dav02} used lensing statistics to constrain the inner slope of lensing galaxies. Using the Schechter function \\citep{sch76}, they constrained the inner slope of lensing galaxies to the range from $1.58$ to $1.98$, at 95\\% confidence level (CL). It is hard to predict how their result would change if the Press-Schechter function \\citep{pre74} were used. Our choice of $\\alpha = 2$ for galaxies is supported by the following fact: stellar dynamics of elliptical galaxies, modeling of lensed systems, and flux ratios of multiple images all give an inner profile that is consistent with SIS \\citep{rix97,rom99,coh01,rus01,tre02,rus02,rus03}.  \\citet{san02} have reported a remarkably flat inner slope in the lensing cluster MS2137-23: $\\alpha < 0.9$ at 99\\% CL. However, by measuring the average gravitational shear profile of six massive clusters of virial masses $\\sim 10^{15} M_\\odot$, \\citet{dah02} have found that the data are well fitted by a mass density profile with $\\alpha \\sim 0.9-1.6$ for SCDM model and $\\alpha \\sim 1.3-1.6$ for LCDM model, both at 68\\% CL. So, our choice of $\\alpha = 1.3$ for population B looks reasonable. The inner density slopes for small mass halos are not well constrained. CDM simulations generally predict a cusped inner density, while other dark matter models, like warm dark matter \\citep{bod01}, repulsive dark matter \\citep{goo00}, and collisional dark matter \\citep{spe00}, tend to predict flatter inner density (see also Ricotti 2002). Our choice of $\\alpha = 1.3$ for population C should be a reasonable upper limit. In her recent paper, by requiring that the Schechter luminosity function and the Press-Schechter mass function to give consistent predictions for the image separation below $1^{\\prime\\prime}$, \\citet{ma03} has shown that the fraction of SIS halos peaks around mass of $10^{12} M_\\odot$ and quickly drops for large and small mass halos. This is qualitatively consistent with the model that we adopt in this paper.  The paper is organized as follows: In \\S\\ref{sec2} we write down the lensing cross-section produced by SIS and GNFW halos. In \\S\\ref{sec3} we show how to calculate the lensing probability, assuming that halos are composed of the population defined above, whose mass function is given by the Press-Schechter function \\citep{pre74}. In \\S\\ref{sec4} we present our results. In \\S\\ref{sec5} we summarize and discuss our results. ",
        "conclusions": "\\label{sec5}  As an extension of our previous work (LO02), we computed the lensing probability produced by a compound population of dark halos. We have calculated the lensing probability for both image separation and time delay. The calculations confirm our previous results (LO02) that the lensing probability produced by GNFW halos with $\\alpha \\la 1.3$ is lower than that produced by SIS halos with same masses by orders of magnitudes, where $-\\alpha$ is the inner slope of the halo mass density. So, for the compound population of halos, both the number of lenses with large image separation ($\\Delta\\theta \\ga 5^{\\prime\\prime}$) and the number of lenses with small image separation ($\\Delta\\theta \\la 10^{-2\\,\\prime\\prime}$) are greatly suppressed. The same conclusion holds also for the number of lenses with large time delay ($\\Delta t \\ga 10\\,{\\rm years}$) and the number of lenses with small time delay ($\\Delta t \\la 10^{-4}\\,{\\rm year}$). (See Figs.~\\ref{fig1} and \\ref{fig5}. This conclusion holds even when the effect of magnification bias is considered, see Figs.~\\ref{fig4a} and \\ref{fig4}.)  We have also tested the dependence of the lensing probability on the redshift of the source object (Figs.~\\ref{fig2}, \\ref{fig3}, \\ref{fig6}, and \\ref{fig7}). The results show that, the lensing probability is quite sensitive to the change in the redshift of the source object. The number of lenses significantly increases as the source redshift increases. However, the rate of the increase decreases as the source redshift becomes large, which is caused by the fact that the proper cosmological distance approaches a finite limit when $z_S \\rightarrow \\infty$. Another interesting result is that, the peak of the lensing probability for each population moves toward large image separation or time delay, as the source redshift increases.  We see that population C (dwarf halos) in an LCDM model has a unique signature in the time domain, c.f. Figures~\\ref{fig5} and \\ref{fig6}. Time delays of less than $10$ seconds and greater than 0.1 second are predicted and should be found in gamma-ray burst sources which are at cosmological distances and have the requisite temporal substructure. Variants of CDM, such as warm dark matter \\citep{bod01}, repulsive dark matter \\citep{goo00}, or collisional dark matter \\citep{spe00} would not produce this feature. However, current surveys do not go deep enough to provide a sufficiently large sample to test the prediction. When more observational data on gamma-ray burst time delay and small splitting angles become available, our calculations can be used to distinguish different dark matter models \\citep{nem01,wil01}.  We have compared the distribution of the number of lenses over image separation predicted by our model with the updated JVAS/CLASS observational data, with the new {\\em WMAP} cosmological parameters (Fig.~\\ref{fig4}). Since the JVAS/CLASS survey is limited to image separation $\\ge 0.3^{\\prime\\prime}$ \\citep{bro02}, we cannot test our predictions for small image separations. However, in the range that is probed by JVAS/CLASS, we see that both the LCDM and OCDM models fit the observation reasonably well and current data do not allow us to distinguish between the two proposed normalizations for the LCDM spectrum, even though these produce predictions that differ by a factor of roughly $1.4$. An explicit search for lenses with image separation between $6^{\\prime\\prime}$ and $15^{\\prime\\prime}$ has found no lenses \\citep{phi01}, which rules out the SIS model for image separation in this range (LO02). This together with our Figure~\\ref{fig4} supports our model of compound population of halos. For separations greater than $10^{\\prime\\prime}$ the differently normalized LCDM models produce significantly different results, thus producing an additional lever to break the degeneracies in the WMAP results (cf. Bridle et al. 2003)  We have also calculated the distribution of the mean time delay vs image separation for the LCDM model (Fig.~\\ref{fig8}). We see that, the compound model fits observations quite well, better than the model of single population of halos \\citep{ogu02}. The compound model predicts a unique feature in the $\\lg\\Delta\\theta - \\lg\\Delta t_{\\rm med}$ plane: there is a ``step'' corresponding to the transition in mass density profile. This can be better tested when more observation data are available.  A controlled survey of lenses with double the sample size of CLASS, perhaps obtainable via SDSS \\citep{yor00,sto02}, should allow one to better distinguish between LCDM variants and perhaps between LCDM models and those based on quintessence \\citep{cal98} rather than a cosmological constant."
    },
    "0212/astro-ph0212126_arXiv.txt": {
        "abstract": "Motivated by the relative lack of neutral hydrogen around Lyman Break Galaxies deduced from recent observations, we investigate the properties of the \\Lya forest around high redshift galaxies. The study is based on improved numerical SPH simulations implementing, in addition to standard processes, a new scheme for multiphase and outflow physics description. Although on large scales our simulations reproduce a number of statistical properties of the IGM (because of the small filling factor of shock-heated gas), they underpredict the \\Lya optical depth decrease inside 1 Mpc~$h^{-1}$ of the galaxies by a factor of $\\approx 2$. We interpret this result  as due to the combined effect of infall occurring along the filaments, which prevents efficient halo gas clearing by the outflow, and the insufficient increase of (collisional) hydrogen ionization produced by the temperature increase inside the hot, outflow-carved bubble. Unless an observational selection bias is present, we speculate that local photoionization could be the only viable explanation to solve the puzzle. ",
        "introduction": "Galaxies form from the intergalactic medium (IGM), process such gas into stars, and possibly re-eject a fraction of it, enriched by nucleosynthetic products, back into the intergalactic space via powerful supernova-driven outflows (Mac Low \\& Ferrara 1999; Ferrara, Pettini \\& Shchekinov 2000; Madau, Ferrara \\& Rees 2001; Scannapieco, Ferrara \\& Madau 2002; Theuns \\etal 2002). The energy deposition connected to these processes is expected to leave at least some detectable imprints on the physical state of the IGM. Thus, it is conceivable that such signatures can be studied through QSO absorption line experiments. Naively, the presence of hot outflowing gas should result primarily in two effects: [i] a decrease of the gas density and [ii] an increase of the temperature caused by shock-heating (acting in conjunction with photo-heating by the UV background) in a large (several hundreds kpc) region around the perturbing galaxy. Both these occurrences would imply an increasingly more transparent Ly$\\alpha$ forest when approaching the galaxy, \\ie a galactic proximity effect. Quantitative confirmation of this scenario has faced tremendous difficulties, standing the complications of the physics of star formation, explosions and metal mixing in multiphase media. Hence, most simulations to date had to rely on {\\it ad hoc} recipes for such processes. \\\\ Nevertheless, these ideas have stimulated the first challenging observations aimed at detecting the imprints of galaxy-IGM interplay. Adelberger \\etal (2002, A02) obtained high resolution spectra of 8 bright QSOs at $3.1<z<4.1$ and spectroscopic redshifts for 431 Lyman-break galaxies (LBGs) at lower redshifts. By comparing the positions of the LBGs with the Ly$\\alpha$ absorption lines in QSO spectra, indeed they conclude that within $\\approx 0.5 h^{-1}$ (comoving) Mpc of the galaxies little \\HI              is present; on the contrary, between 1 and 5$h^{-1}$ Mpc an \\HI excess with respect to the IGM mean is detected. This simple interpretation might be at odd with the results of a VLT/UVES study of the \\Lya forest in the vicinity of the LBG MS1512-cB58 showing the opposite trend (Savaglio \\etal 2002), \\ie an absorption excess close to the galaxy. \\\\ Numerical simulations have also noticeable difficulties reproducing A02 results as discussed by Croft \\etal (2002) and Kollmeier \\etal (2002); however, these studies lack a self-consistent treatment of multiphase gas structure and/or outflow dynamics. Here we revisit A02 results through SPH simulations (Marri \\etal  2003) that implement a new scheme for multiphase hydrodynamics and, more importantly, a physically meaningful outflow treatment; the full description of the code and of the tests made are given in Marri \\& White (2002). We then derive synthetic absorption line spectra along lines of sight randomly traced through the simulation box at $z\\approx 3$ and compare them directly with A02 data to investigate galactic feedback effects on the IGM. ",
        "conclusions": "\\label{sect4} Motivated by the recent observational results of A02, we have studied, with the help of a set of cosmological simulations including star formation in multiphase gas and outflows from galaxies, the effects of galaxy formation/activity on the properties of the surrounding \\Lya forest. Although on  large scales our simulations can reproduce remarkably well a number of statistical properties of the IGM, they fail to predict the observed \\Lya flux increase in regions close to the galaxies themselves. The success can be ascertain to two concomitant effects: (i) outflows preferentially expand in regions of low-density (voids) thus preserving the filaments responsible for the \\Lya absorbing network; (ii) the hot bubbles fill a relatively small fraction of the cosmic volume ($\\approx 14$\\% in our simulations). Support to the first hypothesis emerges also from an inspection of the velocity field in the surrounding of the most massive galaxy in the simulation: the inflow of gas from low density regions is blocked by the bubble expansion, but it proceeds basically unimpeded along the filaments (see Marri \\etal 2003).\\\\ Much more puzzling is instead the reason for the opacity excess (with respect to real data) we see in the simulation in the inner Mpc~$h^{-1}$. Apparently, simulated outflows do not carry sufficient momentum to disperse the density peak created in the vicinity of the galaxy as a leftover of its formation. Also, a large fraction of the outflow energy is used to counteract the infalling gas ram pressure which tends to pile up the gas into the galaxy. The temperature increase close to galaxies amplifies the magnitude of the collisional ionization rate, which becomes larger than the equivalent photoionization rate for $T\\simgt 10^5$~K. In this case the \\HI neutral fraction $x_{\\nHI}$ is independent of gas density and remains at roughly the same level as in the general IGM. Therefore, collisional ionization is not sufficient to balance the opacity increase induced by the density raise and the transmitted flux drops accordingly to the latter. What are the possible alternative explanations for the observed flux trend? The presence of a bias in the data produced by a preferential selection effect of low \\Lya absorption LBGs has already been suggested by Croft \\etal (2002). Another possibility is provided by photoionization, particularly if one recalls the recent results  by Steidel, Pettini \\& Adelberger (2001), who detected flux beyond the Lyman limit (with significant residual flux at $\\lambda <  912$~\\AA) in a composite spectrum of 29 LBGs at $z = 3.4$, a clue of a conspicuous escape probability of ionizing radiation from these objects. As mentioned above, the A02 data would require an optical depth (or equivalently a $x_{\\nHI}$, assuming a prescribed density profile) a factor $\\approx 2-3$ smaller. If this can be achieved with the ionizing flux coming from galaxies must be proved with detailed radiative transfer calculations which are currently ongoing (Maselli \\etal 2003) The first attempt using simplified analytical and/or post-processing techniques to account for this effect (Croft \\etal 2002, Kollmeier \\etal 2002) have yielded so far negative answers, but a fully self-consistent, physically accurate description of the problem is awaited in order to draw a final conclusion.      \\bigskip This work was partially supported by the Research and Training Network `The Physics of the Intergalactic Medium' set up by the European Community under the contract HPRN-CT2000-00126 RG29185. MB thanks P.Petitjean for discussions and hospitality at IAP. We are grateful to S. Bianchi for help with AUTOVP and discussions."
    },
    "0212/astro-ph0212256_arXiv.txt": {
        "abstract": "{ I describe here the performances of a parallel treecode with individual particle timesteps. The code is  based on the Barnes-Hut algorithm and runs cosmological N-body simulations on parallel machines with a distributed memory architecture using the MPI message passing library. For a configuration with a constant number of particles per processor the scalability of the code has been tested up to $P=32$ processors. The average CPU time per processor necessary for solving the gravitational interactions is within $\\sim 10 \\%$ of that expected from the ideal scaling relation. The load balancing efficiency is high ($\\simgt90\\%$) if the processor domains are determined every large timestep according to a weighting scheme which takes into account the total particle computational load within the timestep. ",
        "introduction": "Numerical simulations play a fundamental role for improving the theoretical understanding of structure formation. This approach has received a large impulse from the huge growth of computer technology in the last two decades. Cosmological N-body simulations are now widely used as a fundamental tool in modern cosmology for testing viable theories of  structure formation. A popular approach for solving the gravitational forces of the system is the tree algorithm \\citep{ap85,he87}. The particle distribution of the system is arranged into a hierarchy of cubes and the force on an individual particle is computed by a summation over the multipole expansion of the cubes. An important point in favor of tree codes is that individual timesteps for all of the particles can be implemented easily, this allows a substantial speed-up  of the force evaluation for a clustered distribution.  An important task is the improvement of the dynamic range of the simulations. Large simulation volumes are required for statistical purposes, but at the same time modelling the formation and evolution of each individual galaxy in the simulated volume requires that a realistic simulation should be implemented with $10^8 \\sim 10^9$ particles. This computational task can be efficiently solved if the code is adapted to work on a parallel machine where many processors are linked together with a communication network. This has led a number of authors to parallelize treecodes \\citep{sa91,wa94,du96,da97,li00,sp01,mi02}. In this paper I present a parallel implementation of a multistep treecode based on the Barnes-Hut (1986, BH) algorithm. The code is cosmological and uses the MPI message library. ",
        "conclusions": ""
    },
    "0212/astro-ph0212457_arXiv.txt": {
        "abstract": "We present first results from three-dimensional hydrodynamic simulations of the high redshift formation of dwarf galaxies.  The simulations use an Eulerian adaptive mesh refinement technique to follow the non--equilibrium chemistry of hydrogen and helium with cosmological initial conditions drawn from a popular $\\Lambda$-dominated cold dark matter model.  We include the effects of reionization using a uniform radiation field, a phenomenological description of the effect of star formation and, in a separate simulation, the effects of stellar feedback. The results highlight the effects of stellar feedback and photoionization on the baryon content and star formation of galaxies with virial temperatures of approximately $10^4$K.  Dwarf sized dark matter halos that assemble prior to reionization are able to form stars. Most halos of similar mass that assemble after reionization do not form stars by redshift of three.  Dwarf galaxies that form stars show large variations in their gas content because of stellar feedback and photoionization effects. Baryon-to-dark matter mass ratios are found to lie below the cosmic mean as a result of stellar feedback.  The supposed substructure problem of CDM is critically assessed on the basis of these results. The star formation histories modulated by radiative and stellar feedbacks are discussed. In addition, metallicities of individual objects are shown to be naturally correlated with their mass-to-light ratios as is also evident in the properties of local dwarf galaxies. ",
        "introduction": "\\label{sec:introduction}  Radiative and stellar feedback have long been recognized as an essential ingredient in galaxy formation and they are thought to play an important role in determining the properties of dwarf galaxies. The smallest dark matter halos are thought to be the sites for the formation of the first stars (Couchman and Rees 1986; Tegmark et al. 1997, Haiman and Loeb 1997; Abel et al. 1998; Bromm, Coppi and Larson 1999; Abel, Bryan and Norman 2000, 2001). Whether these smallest objects will allow a significant fraction of the baryons to participate in primordial star formation is unclear because of the effect of complex radiative feedback processes on molecular hydrogen formation (Haiman, Abel and Rees 2000; Ciardi, Ferrara, Abel 2000; Mahacek, Bryan and Abel 2001, Glover and Brand 2001).  However, the early formation of even a few stars in such low-mass objects may significantly affect the evolution of their baryon content due to supernova explosions driving a significant fraction of the gas away from the gravitational potential of the host.    Stellar feedback was first discussed in the context of primordial globular cluster formation by Peebles and Dicke (1968).  In the context of hierarchical structure formation, the possibility that supernovae may remove a significant fraction of the gas present in dwarf galaxies was discussed by Dekel \\& Silk (1986) and by Couchman \\& Rees (1986). Further analytical treatment of such supernova-driven winds was presented by Babul \\& Rees (1992) and Efstathiou (1992).  Star formation and feedback have also been studied in the context of cosmological structure formation simulations. However, the efficiency of supernova energy feedback in driving the gas out of the dark matter potentials was found to depend on the details of the feedback prescription: Katz (1992) and Weinberg, Hernquist \\& Katz (1997) found that the properties of their simulated galaxies were insensitive to the inclusion of star formation and feedback to the simulations. Navarro \\& White (1993) and Navarro \\& Steinmetz (1999, 2000) found that the effect of supernova feedback depends on whether the energy returned to the ISM is in the form of thermal or kinetic energy. In the former case, the energy was quickly radiated away and the effect of feedback was minimal, while in the latter case feedback was efficient in driving gas away from the dark matter potentials and in suppressing star formation. In addition, Mac Low \\& Ferrara (1999) and Navarro \\& Steinmetz (2000) found that the effect of feedback depends on the total mass of the system under consideration, with the low-mass systems being affected much more than the higher-mass systems.  Energy feedback from supernovae is also believed to have a significant effect on the metallicities of dwarf galaxies. The idea of metals being driven away from galaxies is supported by observations of the chemical enrichment of the IGM (e.g. Ellison et al. 2001 and references therein) as well as observations of dwarf galaxies in he Local Group.  Prada \\& Burkert (2002) have summarized a correlation between mass-to-light ratio and metallicity of dwarf spheroidals, which they attribute to metal loss being more efficient for objects with shallower potential wells.  The ability of supernova-driven winds to remove metals from galaxies has been studied by Mac Low \\& Ferrara (1999) and by Aguirre \\etal (2001). Mac Low \\& Ferrara (1999) modeled the effects of repeated supernova explosions from starbursts in dwarf galaxies with hydrodynamic simulations performed in a pre-determined dark matter gravitational potential and with an initial baryon-to-gas mass ratio extrapolated from observations of higher mass systems. Their results suggest that galactic winds are very efficient in removing metals from dwarf galaxies, with the effect being more pronounced for the least massive objects. For larger energy ejection rates (Mac Low \\& Ferrera considered star formation rates only up to about 0.004 $M_\\odot$/year), the gas in the galaxy can be completely blown away (Mori, Ferrara \\& Madau 2002).  Aguirre \\etal (2001) used results from completed N-body + hydrodynamic simulations including star formation which, however, did not directly incorporate the result of feedback, and assumed a series of  ad hoc algorithms to reconstruct the metal enrichment process starting from the star formation history in each simulated object. They found that most galaxies at redshift $\\sim 3$ should exhibit winds driving metals away from their gravitational potentials and enriching the IGM with metals. Note that such winds can now be studied directly using aborption spectra of quasars behind Lyman break galaxies (Adelberger et al. 2002).    An additional factor affecting the baryon contents of dwarf galaxies is photoionization.  The Jeans mass prior to reionization is approximately a factor $10^5$ smaller than after, i.e.  the minimum masses of dark matter halos that can confine baryons are much less at times before the intergalactic medium has been ionized.  Efstathiou (1992) and Babul \\& Rees (1992) discussed analytically the possibility that a photoionizing background might significantly affect the formation of dwarf galaxies and may be the mechanism preventing the collapse of most of the baryonic material into subgalactic objects at high redshifts.  Thoul and Weinberg (1996) showed, using a one--dimensional hydrodynamics code, that photo-ionization affects the collapse of gas into dark halos with velocity dispersions even in excess of $30\\kms$.  Kepner, Babul and Spergel (1997) have looked at the formation of molecular hydrogen in the cores of proto--dwarf galaxies and Kepner et al. (1999) investigated observational consequences of the picture of delayed dwarf galaxy formation. The scenario for the formation of dwarf galaxies outlined by Ferrara and Tolstoy (2000) also takes into account the feedback from photo-ionization. In their model the first stars in observed dwarfs form prior to reionization. Star formation at later epochs originates from new gas infall as the meta--galactic UV background radiation field decreases.  Gnedin (2000) studied the effect of photoionization on the gas fraction of low-mass objects and found a deviation from the universal baryon fraction at low masses.  Many of these papers, as well as Ciardi et al. (2000), did implicitly predict that not all the substructure seen in dark matter--only simulations should also host luminous objects. Images of multiply lensed quasars may be used to constrain the amount of dark substructure in galaxies (Metcalf and Madau 2001). Interestingly, existing data on such systems favors the existence of such dark matter substructure (Metcalf and Zhao 2002, Dalal and Kochanek 2002).  In this paper, we use high-mass resolution, 3-dimensional hydrodynamic AMR simulations to examine the effect of stellar feedback and photoionization on the gas content, star formation histories and metallicities of dwarf galaxies down to a redshift of 3.  Star formation and feedback are included using a phenomenological recipe. Since star formation, metal and energy feedback are followed concurrently with the dark matter and gas evolution, our simulations {\\em do not require invoking any assumptions } on the nature of supernova driven winds (as in Aguirre \\etal (2001)) or on the shape of the dark matter gravitational potentials (as in Mac Low \\& Ferrara (1999)).  In addition, grid based methods are known to be superior to particle techniques when a multi-phase medium needs to be resolved (Croft et al. 2001, Pearce et al. 2000, Ritchie \\& Thomas 2001, and references therein). Here, we will only discuss aspects of the simulation that are expected to be insensitive to the details of the star formation algorithm.   Our paper is organized as follows: In \\S 2 we describe the simulations carried out as well as relevant aspects of our implementation of star formation and feedback. In \\S 3 we discuss the global star formation history in our cube before and after reionization. In \\S 4 we present the effects of feedback and reionization on the baryon contents, star formation histories and metallicities of individual halos.  Finally, we discuss our findings in \\S 5. ",
        "conclusions": "We have presented the first results of a large-scale three-dimensional cosmological hydrodynamic AMR simulation of a cube of comoving size $\\sim 10$ Mpc living in a $\\Lambda$CDM universe with cosmological parameters $(\\Omega_\\Lambda, \\Omega_m, h) = (0.7,0.3,0.67)$.  Star formation was allowed for, according to a parametric prescription creating collisionless particle out of baryonic fluid condensations exceeding the local Jeans mass.  We have compared the results of two runs, differing in the treatment of the stellar particles: in one case the latter did return energy, momentum and metal feedback to the interstellar medium, while in the other no stellar feedback was allowed for. The effects of reionization were studied by subjecting the baryon fluid to a homogeneous, Haardt \\& Madau radiation field of index -1.5, independent of the cosmic star formation.  Over 300 galaxies are produced in our simulation over the mass range $3\\times 10^9 \\le M_{\\rm tot}/M_\\sun \\le 7 \\times 10^{11}$.  The smallest dwarf galaxies in the local group show kinematic masses of $2\\times 10^7\\Ms$ (Mateo 1998 and references therein). However, these observed galaxies may be hosted in much more extended dark matter halos (Stoehr et al 2002).  It seems premature to argue for a significant modification of the cold dark matter scenarios such as to ask for significant annihilation rates (Cen 2001) to accommodate for their existence. The star formation histories of the smallest dwarf galaxies are also essential in understanding the recently proposed substructure problem raised by Moore et al. (1999) and Klypin et al.  (1999).  It has been argued that the systems associated with the extensive dark matter substructure found in N--body simulations of galaxy formation would not be in accord with the observations. This was one of the reasons for the proposal of Spergel and Steinhardt (2000) that the dark matter may be self--interacting which has attracted significant theoretical interest (Burkert 2000; Kochanek and White 2000; Yoshida \\etal 2000, and references therein). The connection between simulated dark matter substructure and observable quantities is a complex one and must depend dramatically on the intricacies of star formation and feedback in the small potential wells of interest (Bullock et al. 2000, Somerville 2002, Benson et al. 2002).  Our results have shown that feedback from stars is a factor which can deplete halos of baryonic gas and significantly reduce their baryon mass fraction. The gas depletion due to feedback results in halos having a baryon - to - dark matter mass ratio which is consistently below the cosmic mean, $\\Omega_b / \\Omega _{DM}$, even in the very high mass objects. In addition, feedback can result to the formation of dark matter halos empty of stars by z=3, and a baryon fraction so low that their observation would be particularly difficult. The existence of such a population of dark dwarf halos could at least partially explain the problem of missing satellites in the Local Group.  We have found that feedback and reionization affect the star formation rate both globally and locally. The slope of the cosmic star formation history curve changes around the epoch of reionization with the effect being much more pronounced when feedback is present. In addition, most halos assembling after reionization do not form stars by $z=3$, an effect which is again enhanced by the presence of feedback. Furthermore, while before reionization most SF activity originates from high-SFR objects, after reionization there is a shift towards lower SFR halos as the source of most cosmic star formation.  Thus, to the extent that this effect is not an artifact of rapid gas depletion due to our high SF efficiency, a significant amount of the total star formation at intermediate redshifts (3-5) might be taking place in low-SFR halos. This effect has been first discussed by Barkana and Loeb (2000) and is here reported in full cosmological hydrodynamical simulations.  Finally, we find that the metallicity of galaxies is correlated with the mass-to-light ratio, increasing for decreasing metallicity. This trend is consistent with observational results of dwarf spheroidals in the Local Group. However, most of the objects in our simulation will merge to make larger galaxies and only a very small fraction will survive as isolated dwarf sized galaxies until today. So the presented comparison again only serves as an illustration  and has as yet little predictive quality for the formation of dwarf galaxies. Nevertheless, we find it very encouraging that these first results show that with yet larger simulations one will be able to make predictions and direct comparisons with properties of local dwarf galaxies."
    },
    "0212/hep-ph0212024_arXiv.txt": {
        "abstract": "{Strange stars (ReSS) calculated from a realistic equation of state (EOS), that incorporate chiral symmetry restoration as well as deconfinement at high density\\cite{d98} show compact objects in the mass radius curve. We compare our calculations of incompressibility for this EOS with that of nuclear matter. One of the nuclear matter EOS has a continuous transition to ud-matter at about five times normal density. Another nuclear matter EOS incorporates density dependent coupling constants. From a look at the consequent velocity of sound, it is found that the transition to ud-matter seems necessary. } ",
        "introduction": "An exciting issue of modern astrophysics is the possible existence of a family of compact stars made entirely of deconfined u,d,s quark matter or ``strange matter\" (SM) and thereby denominated strange stars (SS). They differ from neutron stars, where quarks are confined within neutrons, protons and eventually within other hadrons (hadronic matter stars). The possible existence of SS is a direct consequence of the so called strange matter hypothesis\\cite{bodmer}, according to which the energy per baryon of SM would be less than the lowest energy per baryon found in nuclei, which is about 930 MeV for $Fe^{56}$. Also, the ordinary state of matter, in which quarks are confined within hadrons, is a metastable state. Of course, the hypothesis does not conflict with the existence of atomic nuclei as conglomerates of nucleons, or with the stability of ordinary matter\\cite{farhi,madsen,Bombaci}.  The best observational evidence for the existence of quark stars come from some compact objects, the X-Ray burst sources SAX~J1808.4$-$3658 (the SAX in short) and 4U~1728$-$34, the X-ray pulsar Her X-1 and the superburster 4U~1820$-$30. The first is the most stable pulsating X-ray source known to man as of now. This star is claimed to have a mass $M* \\sim 1.4 ~M_\\odot~ $ and a radius of about 7 kms \\cite{Li99a}. Coupled to this claim are the various other evidences for the existence of SS, such as the possible explanation of the two kHz quasi-periodic oscillations in 4U 1728~-~34 \\cite{Li99b} and the quark-nova explanation for $\\gamma$ ray bursts \\cite{boma}.  The expected behaviour of SS is directly opposite to that of a neutron star as Fig.(\\ref{MR}) shows. The mass of 4U~1728$-$34 is claimed to be less than  1.1 $M_\\odot$ in Li et al. \\cite{Li99b}, which places it much lower in the M-R plot and thus it could be still gaining mass and is not expected to be as stable as the SAX. So for example, there is a clear answer \\cite{singhini} to the question posed by Franco \\cite{fr01}: why are the pulsations of SAX not attenuated, as they are in 4U~1728$-$34 ?  From a basic point of view the equation of state for SM should be calculated solving QCD at finite density. As we know, such a fundamental approach is presently not feasible even if one takes recourse to the large colour philosophy of 't Hooft \\cite{'t Hooft}. A way out was found by  Witten \\cite{wit} when he suggested that one can borrow a phenomenological potential from the meson sector and use it for baryonic matter. Therefore, one has to rely on phenomenological models. In this work, we use different equations of state (EOS) of SM proposed by Dey et al \\cite{d98} using the phenomenological Richardson potential. Other variants are now being proposed, for example the chromo dielectric model calculations of Malheiro et al.\\cite{mal}.  \\begin{figure}[htbp] \\centerline{\\psfig{figure=fig1.eps,height=5cm}}\\vskip .5cm \\caption{\\label{MR} Mass and radius of stable stars with the strange star EOS (left curve) and neutron star EOS (right curve), which are solutions of the Tolman-Oppenheimer-Volkoff (TOV) equations of general relativity. Note that while the self sustained strange star systems can have small masses and radii, the neutron stars have larger radii for smaller masses since they are bound by gravitation alone.} \\end{figure}  \\begin{figure}[htbp] \\centerline{\\psfig{figure=fig2.eps,height=5cm}} \\vskip .5cm \\caption{Strange matter EOS employed by D98 show respective stable points. The solid line for is EOS1, the dotted line for EOS2 and the dashed line for EOS3. All have the minimum at energy per baryon less than that of $Fe^{56}$} \\label{efit} \\end{figure}  Fig.(\\ref{efit}) shows the energy per baryon for the EOS of \\cite{d98}. One of them (eos1, SS1 of \\cite{Li99a}) has a minimum at E/A = 888.8 $MeV$ compared to 930.4 of $Fe^{56}$, i.e., as much as 40 $MeV$ below. The other two have this minimum at 911 $MeV$ and 926 $MeV$, respectively, both less than the normal density of nuclear matter. The pressure at this point is zero and this marks the surface of the star in the implementation of the well known TOV equation. These curves clearly show that the system can fluctuate about this minimum, so that the zero pressure point can vary. ",
        "conclusions": "~~~In our model the density is large even at the surface ($4.586~n_0$ for $EOS1$) so that the quarks in the fermi sea have high momentum on the average. Therefore, they often come close together. It was shown by Bailin and Love that in certain colour channels \\cite{bl}, the short range one gluon exchange spin-spin interaction can produce extra attraction, forming diquarks. In s-state the pair can be formed only in flavour antisymmetric state. The flavour antisymmetric diquark may form in colour symmetric(6) and spin symmetric (triplet) state, or in colour antisymmetric ($\\bar3$) and spin  antisymmetric (singlet) state. The potential strength of the latter is 6 times the former. The N-$\\Delta$ mass splitting is obtainable from the one gluon exchange potential, the delta function part of which is smeared to a Gaussian,  \\be H_{i,j}~=~- 80.51~\\sigma^3(\\lambda_i.\\lambda_j)(S_i.S_j)~e^{-\\sigma^2 r_{ij}^2}. \\label{diq} \\ee  In the above, the strength of this spin dependent part is taken from Dey and Dey(1984) \\cite{dd}.  The pairing energy of ud pairs in the spin singlet and colour $\\bar3$ state is $-3.84~MeV$ \\cite{MNRAS} \\footnote{With this pairing energy superbursts observed from some astrophysical x-ray sources has been explained very well.}. Our model does not predict diquarks that are permanent - since the quarks must have high momentum transfer if they are to interact strongly with a force that is short-range. The formation and breaking of pairs  give rise to endothermic and exothermic processes leading to fast cooling and superbursts respectively.  We also note that the lowering of energy is not very large, as compared  to the approximately 42 $MeV$ energy difference per baryon (930.6 $Fe^{56}$ to $888~MeV$), seen in EOS1. So we  do not expect any drastic change in the incompressibility or the velocity of sound due to diquark pairing, nor do we expect a phase transition to a colour-flavour locked state.  We are grateful to the referee for inviting a comment on quark pairing in the strange matter scenario."
    },
    "0212/astro-ph0212107_arXiv.txt": {
        "abstract": "Radio galaxies provide a means to determine the coordinate distance, the luminosity distance, the dimensionless luminosity distance, or the angular size distance to sources with redshifts as large as two.  Dimensionless coodinate distances for 55 supernovae and 20 radio galaxies are presented and discussed here.  The radio galaxy results are consistent with those obtained using supernovae, suggesting that neither method is plagued by unknown systematic errors. The acceleration parameter q(z) and the expansion rate H(z) or dimensionless expansion rate E(z) can be determined directly from the data without having to make assumptions regarding the nature or evolution of the ``dark energy.''  The expansion rate E(z) can be determined from the first derivative of the dimensionless coordinate distance, $(dy/dz)^{-1}$, and the acceleration parameter can be determined from a combination of the first and second derivatives of the dimensionless coordinate distance. A model-independent determination of E(z) will allow the properties and redshift evolution of the ``dark energy'' to be determined, and a model-independent determination of q(z) will allow the redshift at which the universe transitions from acceleration to deceleration to be determined directly. Determination of E(z) and q(z) may also elucidate possible systematic errors in the determinations of the dimensionless coordinate distances. ",
        "introduction": "The current method used to study the acceleration of the universe is to take the measured $y(z)$ or $D_L(z)$, determine best-fitting global cosmological parameters, and then use these global cosmological parameters to determine the acceleration of the universe as a function of redshift.  In this process, assumptions must be made concerning the nature and redshift evolution of the mass-energy density of the ``dark energy.\"  A direct, empirical determination of the acceleration of the universe as a function of redshift can be determined using the data, without making any assumptions about the nature or evolution of the ``dark energy.'' This can be done using the equation \\begin{equation} -q(z) \\equiv \\ddot{a} a/\\dot{a}^2 = 1~ +~ (1+z) ~(dy/dz)^{-1} (d^2y/dz^2) \\end{equation} valid for k=0; if $k \\neq 0$, another term $[kr(1+z)/(1-kr^2)](dr/dz)$ must be added to the right hand side. Here, $y$ is the dimensionless coordinate distance $y=H_o (a_or)$.  Thus, equation (5) can be used to empirically determine the redshift at which the universe transitions from acceleration to deceleration without requiring assumptions regarding the nature and redshift evolution of the ''dark energy.'' The supernova and radio galaxy data allow a determination of the dimensionless coordinate distance $y$ to each source, at redshift $z$.  These data can then be used to determine $dy/dz$, and $d^2y/dz^2$; these can then be substituted into eq. (5) to determine q(z).  Equation (5) follows from the RW line element and the relation  $(1+z) = a_o/a(t)$. Our measurements of the coordinate distance $(a_or)$ move along the negative direction of $dr$, so eq. (2) with k=0 implies that $a_o~dr = - (1+z)~dt$, or   $(dz/dt) = -a_o^{-1}(1+z)~(dr/dz)^{-1}$. Differentiating $(1+z) = a_o/a(t)$ with respect to time implies that $\\dot{a} = -a_o~(1+z)^{-2} ~(dz/dt)$. Substituting in for $(dz/dt)$, we find $\\dot{a} = (1+z)^{-1} (dr/dz)^{-1}$.  Differentiating again with respect to time, we find $\\ddot{a}=-(1+z)^{-2}~(dz/dt)~(dr/dz)^{-1}~[1+~(1+z)(dr/dz)^{-1}(d^2r/dz^2)]$, which simplifies to eq. (5) using the expressions given here, and the relation $y(z) = H_o (a_or)$. Note, these expressions include the well-known equation (e.g. Peebles 1993; Weinberg 1972) \\begin{equation} H(z) \\equiv \\dot{a}/a = \\sqrt{1-kr^2}~~[d(a_or)/dz]^{-1}~, \\end{equation} or, for $k=0$ and with $H(z) = H_o E(z)$, \\begin{equation} E(z) = (dy/dz)^{-1}~. \\end{equation}   Since eqs. (5), (6), and (7) are derived without any assumptions regarding the mass-energy components of the universe or their redshift evolution, they can be used to directly determine the function $E(z)$, which contains important information on the ``dark energy'' and its redshift evolution, and to determine the dimensionless acceleration parameter $q(z)$ directly from measurements of $y(z)$. The use of the data to directly determine $E(z)$ and $q(z)$ is underway. ",
        "conclusions": "In a spatially flat universe with non-relativistic matter and quintessence, radio galaxies alone indicate with 84\\% confidence that the universe is accelerating in its expansion at the present epoch (Daly \\& Guerra 2002). Results obtained using the Radio Galaxy Method out to redshifts of two are consistent with those obtained using the Supernova Method out to redshifts of about one.  A model-independent way to use the supernova and radio galaxy data to determine the dimensionless expansion rate $E(z)$ and acceleration parameter $q(z)$ is presented and discussed.  A direct determination of $q(z)$ that is independent of assumptions regarding the nature and evolution of the ''dark energy'' would allow the transition redshift from acceleration to deceleration to be determined, and would help to identify any systematic errors that might plague either the radio galaxy or supernova methods.   A direct determination of $E(z)$ would help to quantify the properties and redshift evolution of the ``dark energy'' and to identify potential systematic errors in the methods used to determine $y(z)$."
    },
    "0212/astro-ph0212331_arXiv.txt": {
        "abstract": "We show that solar neutrino experiments set an upper limit of 7.8\\% (7.3\\% including the recent KamLAND measurements) to the fraction of energy that the Sun produces via the CNO fusion cycle, which is an order of magnitude improvement upon the previous limit. New experiments are required to detect CNO neutrinos corresponding to the 1.5\\% of the solar luminosity that the standard solar model predicts is generated by the CNO cycle. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/astro-ph0212088_arXiv.txt": {
        "abstract": "Recently, we have developed and calibrated the Synthetic Field Method to derive the total extinction through disk galaxies. The method is based on the number counts and colors of distant background field galaxies that can be seen through the foreground object; it is the {\\it only} method capable of determining extinction without {\\it a priori} assumptions about the dust properties or its spatial distribution, and has been successfully applied to NGC 4536 and NGC 3664, two late--type galaxies located, respectively, at 16 and 11 Mpc.  Here, we study the applicability of the Synthetic Field Method to $HST$ images of galaxies in the Local Group, and show that background galaxies cannot be easily identified through these nearby objects, even with the best resolution available today. In the case of M~31, each pixel in the $HST$ images contains fifty to one hundred stars, and the background galaxies cannot be seen because of the intrinsic granularity due to strong surface brightness fluctuations. In the LMC, on the other hand, there is only about one star every six linear pixels, and the lack of detectable background galaxies results from a ``secondary'' granularity, introduced by structure in the wings of the point spread function.  The success of the Synthetic Field Method in NGC 4536 and NGC 3664 is a natural consequence of the reduction of the intensity of surface brightness fluctuations with distance.  When the dominant confusion factor is structure in the PSF wings, as is the case of $HST$ images of the LMC, and would be the case of M~31 images obtained with a 10-m diffraction-limited optical telescope, it becomes {\\it in principle} possible to improve the detectability of background galaxies by subtracting the stars in the foreground object. However, a much better characterization of optical PSFs than is currently available would be required for an adequate subtraction of the wings. Given the importance of determining the dust content of Local Group galaxies, efforts should be made in that direction. ",
        "introduction": "\\label{intro}  \\subsection{Some history} \\label{history}  How much dust is there in galaxy disks, and where is it? Reliable answers to these questions have remained elusive, although they bear directly on our understanding of many astronomical problems. Dust has to be properly accounted for to infer luminosity, stellar composition, and mass distribution correctly. Dust could be veiling correlations between global properties of galaxies and affecting our view of their dynamical evolution. It could be hiding significant amounts of luminous baryonic matter, or could trace extremely cold gas, and hence have implications for the dark matter problem. Dust in disks and halos of foreground spirals could be distorting our perception of the high-redshift universe: it could be partly responsible for the apparent decrease of quasi stellar objects (QSOs) at high redshifts (e.g., \\citealp{fall93}), and for the problems encountered in detecting primeval galaxies. And dust in the precursors of present day galaxies could completely change our most recent picture of the star formation history of the universe \\citep{mada96, meur97}. To give just one more example, the Tully-Fisher correlation as a measure of cosmological distance depends on inclination corrections which could be significantly sensitive to dust.  Much work has been devoted to the problem of dust in galaxies, but different authors have often come up with widely differing results, from disks that appear significantly opaque to those basically transparent, and everything in between (see \\citealp{davi95} and references therein). The virtues and disadvantages of different tests of opacity of spiral disks have been discussed at length in Gonz\\'alez et al.\\ (1998, Paper I), and will not be repeated here. In summary, only one method is sensitive to the total amount of dust in disks, independently of its temperature, and does not require that any assumptions be made, neither about the geometry of the background probes, nor about the distribution of the absorbing dust (e.g., uniform screen or a collection of opaque blobs): the comparison of counts of background field galaxies seen through the disk of the foreground galaxy to counts of galaxies in reference fields.  This technique was pioneered by \\citet{shap51}. Before Paper I, however, it had only been applied to the Magellanic Clouds \\citep{shap51, wess61, hodg74, gurw90}. Furthermore, \\citet{wess61} deemed his results as ``inaccurate'', while \\citet{gurw90} qualified theirs as ``tentative.'' The conceptually very simple idea of galaxy counts is, unfortunately, very difficult indeed to implement in practice. The general distribution of disk stars in the foreground galaxy reduces the contrast of background galaxy images, and foreground star clusters and \\hii\\ regions add strongly to the confusion. A first attempt to measure this effect was made by \\citet{macg75}, who used a plot of galaxy counts vs.\\ star counts in the Small Magellanic Cloud to subtract the effects of ``masking by stars'' from those of extinction. In Paper I, we developed the ``Synthetic Field Method'' (SFM) in order to quantify and calibrate the effects of crowding and confusion, and to determine the accuracy with which statements can be made about the total opacity and reddening of an average line of sight through a foreground galaxy. We then applied it to NGC~4536, a bright spiral galaxy in the direction of Virgo, at a distance of 16.2 Mpc \\citep{saha96a}, and to NGC~3446, a small Magellanic irregular at a distance of  11.9 Mpc. In the SFM, images of a suitable reference field are added directly to the images of the foreground galaxy under study. The extinction through that foreground galaxy is then determined by comparing the numbers of ``real'' background galaxies, not to the number of galaxies in the original reference field image, but to the number of galaxies in the reference field that can still be identified after its addition to the foreground galaxy image. This procedure was hinted at by \\citet{wess61}, and by \\citet{gurw90}; they superimposed negatives of fields with and without foreground galaxy as an experiment to estimate the reduction in galaxy counts owing to crowding and confusion, but unfortunately did not correct their results accordingly.  The advent of the Hubble Space Telescope ({\\sl HST}) has made the method more easily applicable, thanks to its superb resolution and sensitivity, as well as to the stability of the sky, the photometry, and the point spread function (PSF). Other factors that favor the application of the method nowadays are the availability of deep and medium deep {\\sl HST} images that can be used as reference fields \\citep*[Williams et al.\\ 1996; Griffiths et al.\\ 1994; Ostrander et al.\\ 1998][]{ratn99a,ratn99b}, and the access to automated surface photometry packages \\citep[for example,][]{bert96}, which allow an accurate assessment of galaxy detection limits.  De Vaucouleurs (1995) suggested that, after the Magellanic Clouds, the best targets would be the next largest galaxies in angular size, i.e., M~31, M~33, M~51, M~81, M~101. Of course, the surface density of background field galaxies is roughly constant in any direction of the sky, but closer galaxies will cover a larger angular area, allowing for results with better spatial resolution. Indeed, M~31 appears as an excellent potential candidate to apply the SFM because it is so nearby, and so well studied. Not only it could offer a unique opportunity to study the distribution of dust on small scales, and in particular to investigate the differences between arm and interarm regions, and between different galactic radii. But also, large--scale, high--resolution maps of the 21--cm line of \\hi\\ \\citep{brin84}, and of the 2.7--mm line of CO \\citep{loin99,niet2000} have been obtained, as well as large--scale maps of the distribution of warm dust (Haas et al. 1998, and references therein). Once the total dust content is known, it would become possible to determine unambiguously whether there exists a significant fraction of very cold dust (at T $<$ 10 K) -- as has been suggested in the past -- that has so far escaped detection in the far--infrared. Also, assuming a canonical dust--to--gas ratio, the total gas content could be measured by means completely independent from the usual methods. In particular, the mass of molecular gas could be estimated without having to rely on the (highly uncertain -- see Loinard \\& Allen 1998) CO to H$_2$ ``conversion factor''.  \\subsection{It is much harder than it looks} \\label{motiv}  In a nutshell, the motivation for this paper is the fact that -- as we shall see momentarily -- applying the SFM to {\\sl HST} Wide Field Planetary Camera 2 (WFPC2; \\citealp{holt95}) data of M~31 or the LMC has not proven easy or very fruitful. Almost no background galaxies were visible through the disk of M~31, and very few could be seen through the Magellanic Clouds (see below). In the case of the LMC, we were particularly surprised because a very simple calculation (dividing 10$^{10}$ stars by the area of the galaxy or $\\sim$ 68 kpc$^2$) shows that there is a star only every $\\sim$ 4 linear Wide Field Camera (WFC) pixels; i.e., with 0\\farcs1 pixels, we should already be looking between the stars!\\footnote{M~31 is about 14 times further and has more stars per unit area (see \\S\\ref{starsim}) than the LMC, so the WFC Andromeda data should not have ``empty'' pixels. But Virgo galaxies are more than 20 times further away than M~31 and hence have many more stars per pixel, and yet the SFM was successfully applied there.} Moreover, since it is also very difficult to see any reference galaxies in the synthetic combined images, it is impossible to tell whether the paucity of ``real'' background galaxies is caused by extinction or by lack of contrast against the galaxy foreground. In order to illustrate the problem, figure \\ref{lmc_color} shows color cutouts of one of the WFPC2 subfields of the LMC we have used, without ({\\it left}) and with ({\\it right}) the HDF-N added in; finding galaxies, especially the simulated ones, behind NGC~4536 in Virgo, in data with similar surface brightness, had been comparatively much easier -- albeit easy, it was not.  In summary, while the SFM provides useful results at the distance of Virgo, it seems to fail in nearby objects. The purpose of the present paper is to determine quantitatively the conditions in which the SFM can yield useful results, and whether we can expect (or when we can expect) to use it to determine accurately the total opacity of Local Group and nearby galaxies. To this end, we will first look into the characteristics, both of the NGC~4536 images where the SFM was applied successfully, and of those images where it failed. Subsequently, we will perform simulations at different levels. First, we will artificially reproduce the stellar fields of our selected images; we will constrain these simulations with published luminosity functions, as well as with the stellar counts, surface brightnesses, and pixel--to--pixel dispersions of the data. Next, we will simulate either changes of surface brightness at fixed resolution, or changes of resolution (to mimic improvements that can be expected with increasingly bigger, diffraction--limited telescopes). To each simulated image with different brightness/resolution, we will add the HDF-N (used as reference field), in order to estimate the number of background galaxies that we can expect to see in the absence of extinction. ",
        "conclusions": "The analysis of the simulations presented in this paper shows that, for the time being, Virgo seems to offer the best prospects for the application of the SFM. Reasonably accurate statements can be made about the opacity of the disks at that distance (Paper I), and long integration times are not required. Arcsecond resolution \\hi\\ VLA observations could be obtained for a significant number of the spirals in Virgo, and the Atacama Large Millimeter Array (ALMA) will be able to provide high--sensitivity CO maps with the same resolution. Obtaining the opacity of a large sample of spiral in Virgo using the SFM could, therefore, be used to statistically investigate the properties of the ISM there. Objects more distant than Virgo would tend to present even smoother foreground stellar fields, but the opacity there would be determined with less accuracy because their smaller angular size would reduce the number of background galaxies that could be identified.  For Local Group galaxies like M~31 and the LMC, the SFM will hardly be usable, even with telescopes significantly better than those currently at hand, unless the PSF is better understood. An adequate subtraction of individual point sources from WFPC2 data is an extremely difficult task, owing both to the undersampling \\citep{ande00}, and to the variability of the PSF with time, position in the chip, and spectral energy distribution of the imaged objects; the variability affects especially our present ability to model the wings \\citep{kris01}, i.e., the source of the secondary granularity that hampers the detection of the background galaxies. Since the quality and stability of the PSFs of future instruments are unlikely to be much better than those of the WFPC2 PSF, a detailed understanding of optical/near-IR PSFs will be required to improve the results of the SFM in any Local Group galaxy."
    },
    "0212/astro-ph0212041_arXiv.txt": {
        "abstract": "Gravitational microlensing has  proven to be a powerful  probe of both the structure at the heart of quasars and the mass function of compact objects in  foreground lenses.  This  paper examines the  potential of gravitational  microlensing  in  probing  the scale  of  structure  in absorbing material  within the lensing  galaxy. We find that,  in this high optical  depth regime,  significant variations in  the equivalent width of absorption  features can be induced, although  the details of these  are dependent  upon the  scale  of structure  of the  absorbing material.  The  paper concludes with an examination  of the absorption line  variability  observed   in  the  gravitationally  lensed  quasar PKS1830-211, demonstrating how this may indicate the presence of small scale structure in  the cold molecular gas present  within the lensing galaxy. ",
        "introduction": "With  the  advent  of  dedicated  monitoring  programs,  gravitational microlensing of quasars  has provided some of the  strictest limits on the scale  of the central accretion  disk (e.g. Wyithe,  Agol \\& Fluke 2002) and  extended emission  line material  (Lewis, Irwin,  Hewett \\& Foltz 1998).  As well as  imaging the source, microlensing also probes the nature of the lensing  galaxy, with the statistics of microlensing variability constraining the mass  function of the microlensing masses in the intervening system (Wyithe et al. 2000; Wyithe \\& Turner 2001).  In this paper we demonstrate how microlensing provides a further probe of the distribution  of material within a lensing  galaxy. Rather than focusing  upon  the point-like  microlensing  masses,  we examine  the influence of material  responsible for producing absorption signatures in the spectrum of the more distant quasar.  This absorption signature is time dependent, as relative  movement of the microlenses results in differing microimage brightnesses.  In Section~\\ref{microsection} we outline the general approach, with an extension  into   the  high  optical  depth   microlensing  regime  in Section~\\ref{high}.  Sections~\\ref{clouds}  and \\ref{magmaps} describe the cloud  distributions employed in this study  and the magnification maps  from which  the results  are taken.   Section~\\ref{pks} examines whether the  radio spectral variability  observed in the  lensed radio quasar  \\pks\\ is  consistent  with the  hypothesis  presented in  this paper,  while   the  conclusions  to  this  study   are  presented  in Section~\\ref{conclusions}. ",
        "conclusions": "Employing  a  numerical  approach,  this paper  has  investigated  the influence  of a  combination  of microlensing  and  a distribution  of absorbing material within the  lensing galaxy upon temporal changes in the depth of absorption lines in quasar spectra.  It is found that the action  of macrolensing  compresses the  large scale  absorption cloud information into  smaller region, as  seen by the source,  whereas the microlensing   `folds'   the   absorption  patten.    Several   clouds distributions were  examined, and significant modulation  of the depth of  the absorption  line resulted.   The form  of the  absorption line variability  differed significantly  from  the simple  case where  the light from the quasar shone solely through an absorbing medium with no additional   influence  from   gravitational  lensing.    Hence,  such variability   could    influence   systems   suffering   gravitational microlensing. Due to the small  size of the Einstein radius of stellar mass objects at cosmological distances, however, only sub-parsec scale variations in absorbing material will result in an observable effect.  \\begin{figure} \\centerline{ \\psfig{figure=fig15.ps,angle=270,width=3.3in}} \\caption{As Figure~\\ref{Temporal}, but for models Run2C and Run2D.} \\label{Temporalc} \\end{figure}  The  gravitationally  lensed quasar,  \\pks,  possess several  temporal features  that   are  consistent  with   it  undergoing  gravitational microlensing.   Furthermore,  \\pks\\ possesses  a  number of  prominent molecular  absorption lines visible  at radio  wavelengths. In  one of these,   ${\\rm  HCO^+(2  \\leftarrow   1)}$,  small,   but  significant variations in the absorption line  profile over a period of six months have been reported. The analysis presented in this paper suggests that these  changes   may  be  due   to  the  influence   of  gravitational microlensing.  It  is important to note, however,  that currently very little is  known about  the scale of  structure of  absorbing material within  external galaxies  and hence  the identification  of  any line variability  as being due  to microlensing  may be  problematic. While this paper has demonstrated  that microlensing induced absorption line variability is  not perferctly correlated with the  variability of the brightness  of the  macrolensed images,  some correlation  does exist. Hence,  long  term photo-spectral  monitoring  at  high resolution  of potential systems, such as \\pks,  is required before such a conclusion is confirmed.  \\begin{figure} \\centerline{ \\psfig{figure=fig16.ps,angle=270,width=3.4in}} \\caption[]{\\label{SpecAbs} These  three panels present  the absorption variation (lighter line) and light  curves (darker line) for the lower path of Run1A (Fig.~\\ref{run1a}). The smallest source is considered in the upper panel, with the largest in the lower panel. It is clear that while there  is some correlation  with the absorption  variability and the light curve variations, this does not always hold.  } \\end{figure}"
    },
    "0212/astro-ph0212277_arXiv.txt": {
        "abstract": "We present continuum observations of the edge-on spiral galaxy NGC 5775 at 617 MHz with the Giant Metrewave Radio Telescope (GMRT) and at 850 $\\mu$m with the Submillimetre Common-User Bolometer Array (SCUBA) on the James Clerk Maxwell Telescope (JCMT).  We report the detection of dust at high-latitudes, extending to about 5 kpc from the disk of the galaxy, and spurs of non-thermal radio emission. The 617 MHz and 850 $\\mu$m  distributions are compared with one another and we find a strong correlation between metre wavelength radio emission and sub-mm emission in the disk and at high-latitudes. This suggests that there may be a more fundamental relationship between cold dust and synchrotron radiation other than the link via star formation. ",
        "introduction": "There is now considerable evidence for galactic halos, thick disks, and discrete connecting features between the disk and the halo in many star-forming spiral galaxies.  Such features are best detected directly by observing edge-on galaxies with sufficient sensitivity and spatial resolution.  High-latitude (i.e. $\\geq$ 1 kpc) emission has now been observed in every interstellar medium (ISM) tracer (cf. Dahlem 1997). NGC 5775 is an example of a galaxy in which every ISM component has been observed at high-latitudes (Lee et al. 2001).  The Infrared Astronomical Satellite (IRAS) data of a sample of spiral galaxies showed the estimated ratio of their gas mass to dust mass, to be about 10 times larger than the average value for the inner region of the Galaxy (Devereux \\& Young 1990). It was suggested that this could be due to a large amount of cold dust in these galaxies that IRAS was not sensitive to.  Observations with good spectral coverage and sensitivity were required to disentangle the warm and cold dust components. The Submillimetre Common-User Bolometer Array (SCUBA) operating on the James Clerk Maxwell Telescope (JCMT), which we have used for this project, allows astronomers to determine the distribution and quantity of such cold dust with high resolution. Currently, only  a small number of quiescent and active edge-on spiral galaxies have had their dust distributions mapped using either SCUBA or the IRAM 30-m telescope; these include NGC 891 (Alton et al. 1998), NGC 3079 (Stevens \\& Gear 2000), NGC 4631 (Neininger \\& Dumke 1999), NGC 4565 (Neininger et al. 1996), NGC 5907 (Dumke et al. 1997) M82, NGC 253, NGC 4631 (Alton et al. 1999) and now NGC 5775.  In the 1970s a correlation was discovered between mid-infrared ($10\\mu$m) and radio (1425 MHz) luminosities from the nuclei of Seyfert galaxies, and later normal spiral galaxies (van der Kruit 1971, 1973; Rieke 1978).  In the 1980s the IRAS mission established a universal and tight correlation between the far infrared (FIR) and the radio continuum luminosities of galaxies, (de Jong et al. 1985; Helou et al. 1985).  Recently a number of authors have investigated the FIR/radio correlation within the disks of nearby galaxies in an attempt to gather more detailed information on how the correlation works (Beck \\& Golla 1988; Bicay \\& Helou 1990, amongst others).  Presently, it is widely believed that the main reason behind the FIR/radio correlation, is that both emissions depend on the same recent star formation activity of that particular galaxy (Condon 1992). That is, hot young stars heat the dust, producing FIR emission, as well as ionize gas and provide a supply of SNe, producing the radio continuum emission.  However, the situation may not be so clear cut. For example, in a detailed study of M31, Hoernes, Berkhuijsen \\& Xu (1998, hereinafter referred to as HBX), reported the discovery of a nonthermal-radio/cool-dust correlation, quite separate from the thermal-radio/warm-dust correlation that is linked through massive ionizing stars. It is worth noting that HBX did not find evidence of a correlation between either non-thermal radio with warm dust, or thermal radio with cool dust. Thermal emission from warm dust is the result of heating from massive stars ($>20 M_{\\odot}$), while cool dust could be powered by intermediate mass stars (5-20 $M_{\\odot}$) and the interstellar radiation field. The nonthermal-radio/cool-dust correlation could not be explained through a dependence on the same energy source.  HBX suggest a possible scenario involving coupling of gas, cold dust and magnetic fields.  In this paper we present the first detection of high-latitude dust for NGC 5775.  Observing with SCUBA, we map the cold dust in the disk and halo. We also present radio continuum observations of the galaxy taken with the Giant Metrewave Radio Telescope (GMRT), which trace largely the non-thermal synchrotron emission. The observations from SCUBA and GMRT are spatially compared with one another in an effort to investigate whether a cold dust/non-thermal radio emission correlation exists in the disk and halo of this galaxy. NGC 5775 is an edge-on spiral galaxy at a distance of 24.8 Mpc (H$_o$=75 km s$^{-1}$ Mpc$^{-1}$), and with an inclination angle of 86$^\\circ$ (Irwin 1994).  NGC 5775 is found on the FIR/radio correlation at log P$_{20 cm}$ = 22.32 W Hz$^{-1}$ and log (L$_{FIR}$/L$_\\odot$) = 10.44. The results in this paper are the first ones of an ongoing project to investigate a possible sub-mm/metre wavelength correlation in the disks and halos of galaxies.  \\section[]{Observations and Data Reduction}  \\subsection{Submillimetre} Observations of NGC 5775 were made on 2001 January 4 and 5, at the JCMT using SCUBA.  SCUBA images a $2.3^{\\prime}$ region of sky simultaneously at 450 and 850 $\\mu$m. Our observations provide fully sampled images by moving the secondary mirror in the 64pt jiggle map mode.  We performed pointing checks every hour against a bright source. Skydip measurements were also made to determine the atmospheric transparency by change in elevation (opacity was about 0.27 at 850 $\\mu$m and 1.67 at 450 $\\mu$m).  Mars was mapped twice each night for calibration of our source and to determine the beam shape. Our observations of NGC 5775 consisted of three exposures of 45 minutes each over two overlapping fields, since our galaxy is larger than the diameter of the region that the telescope can image.  The software package SURF was used to reduce the SCUBA data.  All maps were reduced using typical reduction tasks.  SURF was used to clean and flat-field the images.  Dirty bolometers were removed and the image noise was reduced using a SURF task which compensated for spatially correlated emission. The pointing drifts were corrected and the images were calibrated using Mars. Finally all the observations were combined into a single map.  The resulting map has a resolution of 14.6$^{\\prime\\prime}$ and an rms noise of 4 mJy/beam.  \\subsection{Radio Continuum} NGC 5775 was observed with the GMRT on 2000 July 1 and 16 at 617 MHz using a bandwidth of 16 MHz.  The GMRT consists of thirty steerable, 45-metre diameter antennas, of which 14 are located randomly in a central 1 km $\\times$ 1 km diameter region referred to as the Central Square.  The remaining antennas lie along 3 arms of an approximately `Y' configuration, with each arm extending to a length of about 14 km.  This provides $uv$ coverage for imaging both the compact and extended emission.  A more detailed description of the GMRT can be found in Swarup et al. (1991).  These observations were made during the commissioning phase of the GMRT and typically 25 to 27 antennas were available for each day.  The source was observed on two different days for checking consistency of the images.  The primary flux density calibrator was 3C286 while the phase calibrator was 3C298.  The total time spent on the source was approximately 120 minutes on each observing day.  The analysis of the data was done using the Astronomical Image Processing System (AIPS) from the National Radio Astronomy Observatory.  The data editing itself was done within the AIPS environment using the task UVFLG. However, the identification of bad data was done outside the AIPS environment with algorithms developed by J. Stil in a package called `borg'. The borg allows the user to define what constitutes bad data. For our data, borg identified all points that had near zero flux density at each channel for the flux density and phase calibrator. The routine also identified the spikes in the data while observing the calibrators by finding points that deviate more than, approximately 3 sigma from the modal amplitude for a particular baseline, channel and source. If any baseline had a significant number of points flagged, it was subsequently flagged from all sources. The borg package helped to ease the editing of the data because the GMRT did not have any automatic flagging of bad data, including bad antennas at the time of these observations.  The AIPS tasks CALIB, CLCAL and BPASS were used on the flagged data set to calibrate the source.  The central 100 channels of the total of 128 channels were used in the analysis. The channels were averaged and  Fourier inversion and cleaning were performed on the NGC 5775 data using IMAGR. A series of self-calibration and cleaning procedures were then conducted to produce two maps, one from each day of observations.  After determining that the two maps were virtually identical we used the AIPS task DBCON to combine the UV plane data from each observation and produced a final single map to increase our signal-to-noise ratio. Using the brightest source visible in our observation and the rms noise of the map, 0.9 mJy/beam, we calculate a dynamic range of 112.  Although the angular resolution using the entire GMRT is approximately 6 arcsec, we have tapered the GMRT images to 14.6$^{\\prime\\prime}$ to be comparable with the 850 $\\mu$m SCUBA map. ",
        "conclusions": "\\subsection{Correlation} The similarities between the sub-mm and radio maps demonstrate that a correlation exists between the radio continuum and sub-mm radiation in NGC 5775. A more quantitative demonstration is shown in Fig. 7, where we have plotted the flux densities at 850 $\\mu$m against those at 617 MHz at identical independent positions for the entire disk and halo of NGC 5775. Fig. 8 shows the correlation for positions solely in the disk.  It is clear that sub-mm and radio emission are highly correlated. The slope of the linear-least squares fit to the entire disk and halo is 0.26$\\pm$0.01 (Fig. 7), while for the disk alone the slope is $0.25\\pm0.03$ (Fig. 8). The corresponding values for the north-eastern and south-western halos are $0.31\\pm0.04$ and $0.26\\pm0.03$ respectively (Fig. 9). Within the errors, the slopes for each section of the galaxy are similar, suggesting that the reason behind this correlation for different parts of the galaxy must also be similar.  This result also suggests that the discrete kpc disk halo features originate in the disk of the galaxy.  The fainter dotted line in Fig. 8 is the best fit of a power law regression, which is similar to the linear regression.  The slope of the power law is $0.85 \\pm 0.08$.  This value is very close to one, indicating the relationship is approximately linear.  For comparison, the slopes typically found for the radio-FIR correlation are between 1.1 and 1.3.  The sub-mm observations map the cold dust, which emits predominantly longward of a couple of hundred $\\mu$m (Sauvage 1997). The detection of cold dust was not a surprising finding since dust-to-gas mass ratio of external galaxies derived from IRAS were approximately 10 times lower than in our galaxy (cf. Devereux \\& Young 1990).  The IRAS observations are largely sensitive to warm dust and helped establish the tight and universal correlation between FIR and radio continuum emission. While this correlation is attributed to the dependence of both emission processes on recent star formation activity, the 617 MHz/850 $\\mu$m correlation demonstrated here is unlikely to be the result of a common energy source for the emission at the two wavebands. The cold dust in galaxies is heated not by massive ionizing or intermediate mass non-ionizing stars, but possibly by the general interstellar radiation field (Hippelein et al. 2000).  Also, the radio map at 617 MHz, is almost entirely due to non-thermal synchrotron radiation.  From Condon (1992) we estimate the ratio of thermal flux density to total flux density at 617 MHz in NGC 5775 to be approximately 0.06 using a spectral index of 0.93 (Irwin et al. 1999).  The ratio of the thermal component is approximately twice the above value at 1420 MHz.  \\begin{figure} \\centering\\epsfig{figure=fig7.eps, height=9cm} \\caption {The 850 $\\mu$m flux density plotted against 617 MHz flux density for the disk and halo of NGC 5775. The line indicates the linear least-squares fit to the data.} \\end{figure}  Further supporting evidence that the correlation is not due to a simple shared energy source, comes from the halo region of both maps.  The overlay in Fig. 4 demonstrates that there is a spatial relation between the sub-mm and metre-wave radio emission in the halo region of the galaxy.  This relation is more clearly shown in Fig. 9, which is similar to Fig. 7 except that only high-latitude emission ($>$ 1.5 kpc) is shown.  The circles represent positions located in the halo on the north-eastern side of the galaxy axis, while the plus signs mark positions in the south-western halo. In these regions there is no evidence of significant star formation. The link between the two types of emission must be due to another factor.  \\begin{figure} \\centering\\epsfig{figure=fig8.eps, height=9cm} \\caption{The 850 $\\mu$m flux density plotted against 617 MHz flux density for the disk of NGC 5775. The solid line indicates the linear least-squares fit to the data. The dotted line shows the best fit for a power-law relationship.} \\end{figure}  \\begin{figure} \\centering\\epsfig{figure=fig9.eps, height=9cm} \\caption{The 850 $\\mu$m flux density plotted against 617 MHz flux density for the halo region. The halo on the north-east side of the galaxy axis of NGC 5775 are represented by circles, and the linear least-squares fit by the dotted line. The halo on the south-west side of the galaxy axis is represented by the plus sign, and the least-squares fit by the solid line.} \\end{figure}  Our findings lead us to suggest that the well-known FIR/radio correlation might be one aspect of a more fundamental relationship.  These results on NGC 5775 not only demonstrate a relationship between the two types of emission, namely thermal emission from cold dust and non-thermal emission from relativistic electrons, but could provide clues towards understanding the distributions and relationships between gas, dust and magnetic fields.  Studying the more abundant cold dust component, which could contribute approximately 80 to 95\\% of the dust in normal spiral galaxies (Sievers et al. 1994), has revealed a correlation between the non-thermal radio and sub-mm emission in NGC 5775. This could indicate a wider and more important relationship than one based on recent star-formation activity alone.  Currently a model that can perhaps explain the correlation presented in this paper is the magnetic field-gas coupling model described by HBX to explain the relationship between the non-thermal and cold-dust emission in M31.  In this model the cold dust and the non-thermal radio continuum emission are linked through the magnetic field coupled to the neutral gas which is mixed with the cold dust.  Using this model, HBX could explain the slope of $0.80 \\pm 0.09$ for the observed relationship between the non-thermal radio and cool dust emission in M31 (HBX). This slope is very close to our value of $0.85 \\pm 0.08$ for NGC 5775.  An investigation of the magnetic field in NGC 5775 was conducted by T\\\"{u}llmann et al. (2000).  An analysis of polarized radio continuum emission at 1.49 and 4.86 GHz indicated that the magnetic field in the halo has a strong component that is aligned with radio continuum spurs and are in a direction perpendicular to the disk of NGC 5775.  The magnetic fields extend to sufficient latitudes from the disk and could provide an explanation of the extra-planar features presented above at 850 $\\mu$m and 617 MHz. This supports the possibility that a magnetic field-gas coupling model may explain the cold dust non-thermal gas correlation we have found in NGC 5775.  It is interesting to note that NGC 5775 may be interacting with the nearby galaxy NGC 5774 (Lee et al. 2001), and enquire whether the dust which has been detected up to $\\approx$5 kpc feeds the intergalactic medium. However, without an estimate of the the velocity of the dust we cannot be certain that these dust grains will escape NGC 5775's gravitational well.  However, a few of the high-latitude features appear to suggest that dust might be falling back towards the disk. This is interesting in light of observations by Lisenfeld et al. (2002) which show no dust in the halo of the dwarf galaxy NGC 1569.  The mass difference between dwarf galaxies and normal spirals may account for the missing dust, if the dust escaped into the intergalactic medium."
    },
    "0212/nucl-th0212096_arXiv.txt": {
        "abstract": "We explore the relevance of color superconductivity inside a possible quark matter core for the bulk properties of neutron stars.  For the quark phase we use a Nambu--Jona-Lasinio (NJL) type model, extended to include diquark condensates. For the hadronic phase, a microscopic many-body model is adopted, with and without strangeness content. In our calculations, a sharp boundary is assumed between the hadronic and the quark phases. For NJL model parameters fitted to vacuum properties we find that no star with a pure quark core does exist. Nevertheless the presence of color superconducting phases can lower the neutron star maximum mass substantially. In some cases, the transition to quark matter occurs only if color superconductivity is present. Once the quark phase is introduced, the value of the maximum mass stays in any case below the value of two solar masses. ",
        "introduction": "\\label{1} According to QCD, the fundamental theory of strong interaction, nuclear matter is viewed as the confined phase of quark-gluon matter, where chiral symmetry is broken and quarks are bound inside nucleons. At large enough density, nuclear matter is expected to undergo a phase transition to the deconfined phase, where quarks and gluons are free to move in the medium and chiral symmetry is restored. Unfortunately, this transition is not well understood, since QCD lattice calculations cannot yet be performed at large density, i.e., at large chemical potential. Experimentally, in ultra-relativistic heavy ion collisions the transition to the deconfined phase is expected to occur at high temperature but essentially at zero baryon density. On the other hand, neutron stars are believed to contain very high baryon density in their interior, where the transition to the deconfined phase could occur. It has been argued by several authors that the bulk properties of neutron stars can be strongly affected by the presence of a core where a quark phase or a mixed hadron-quark phase is present. The fact that the quark phase - if present - is likely to be a color superconductor has recently attracted much attention (For reviews see \\cite{RaWi00,Alford01} and references therein). In general, the condensation energy associated with the superconductivity is only a small fraction of the total energy, and the bulk properties of the equation of state (EOS) are hardly affected. This is because only a small fraction of fermions around the Fermi surface participates effectively to pairing, and the condensation energy is not proportional to the pairing gap $\\Delta$, but to $\\Delta^2/E_F$, where $E_F$ is the Fermi energy. These considerations equally apply to color superconductivity in quark matter. However, as recently argued~\\cite{AlRe02}, the structure of neutron stars appears to be sensitive to the properties of the possible quark phase and even a relatively small change in the quark equation of state could affect the predictions of neutron star properties. \\par In this paper we study the consequence of introducing explicitly color superconducting phases in the deconfined quark matter for the structure and bulk properties of neutron stars. To this purpose, we describe the hadronic phase by the EOS derived from a microscopic many-body theory \\cite{bbb}. For the quark matter we adopt a three-flavor Nambu-Jona Lasinio (NJL) type model, which has been recently extended to include diquark condensates~\\cite{BO02,SRP02,NBO02}. In contrast to bag models which are often employed to describe quark phases, the NJL model dynamically generates a density-dependent effective bag constant and density dependent effective quark masses which can be considerably larger than the bare masses. It has been shown earlier that this has important consequences for the (non-) existence of absolutely stable strange quark matter \\cite{BO99} and the color-flavor (un-) locking phase transition in quark matter~\\cite{BO02}. In the context of neutron star interiors the NJL model has been employed in Refs.~\\cite{SLS99} and \\cite{SPL00}. It was found by both groups that quarks exist at most in a small mixed phase regime but not in a pure quark core. This could be traced back to the relatively large values for effective bag constant and the effective strange quark mass which results from the calculation. On the other hand these calculations did not include diquark condensates. As recently discussed within a bag model~\\cite{AlRe02} the latter effectively lower the bag constant and could have sizable consequences for the hadron-quark phase transition and for the properties of compact stars. It is therefore interesting to study these effects for an NJL model equation of state.  Depending on the flavors which participate in a diquark condensate one can distinguish between several color superconducting quark phases, most important the two-flavor superconducting (2SC) phase~\\cite{ARW98,RSSV98} where only up and down quarks are paired and the color-flavor locked (CFL) phase~\\cite{ARW99} which contains $ud$, $us$, and $ds$ pairs.  In principle this can give rise to a large number of globally electric and color neutral mixed quark phases which are, however, unlikely to be stable if Coulomb and surface effects are included~\\cite{NBO02}. Therefore in this letter we discuss only homogeneous neutral quark phases which have been constructed first in Ref.~\\cite{SRP02}. We have checked, however, that the use of the mixed quark phase equation of state of Ref.~\\cite{NBO02} yields practically the same results.  Similarly, we assume a sharp (first order) transition from the hadronic to the quark phase.  This is motivated by the results of Ref.~\\cite{ARRW01} where it was found that a quark-hadron mixed phase is unlikely to be stable for reasonable values of the surface tension. We will briefly comment on the possible consequences of a mixed phase in Sec.~\\ref{stars}. ",
        "conclusions": "We have combined various hadronic EOS with an NJL model EOS with and without color superconductivity to construct the hadron-quark phase transition and to investigate the structure of compact stars. Our results do not allow for a pure quark phase in the interior of a neutron star. In some cases we find no phase transition at all while in other cases the star becomes unstable as soon as the phase transition occurs. Both phenomena are mainly due to the relatively large values of the effective strange quark mass and the effective bag constant in the NJL model. Considering non-superconducting quark matter we thus confirm the results of Ref.~\\cite{SLS99}. Including color superconductivity does not change these findings.  Of course, the properties of the NJL EOS depend on model parameters. In this letter we have adopted the parameters of Ref.~\\cite{Rehberg} which have been adjusted to vacuum properties. Although this is common practice in this field, it is very well possible that these parameters are not appropriate to describe quark matter at densities relevant for compact stars. Thus our arguments could also be turned around: If there were strong hints for the existence of pure quark cores in compact stars, this would indicate a considerable modification of the effective NJL-type quark interactions in dense matter~\\footnote{After submission of our manuscript, two groups have reported the construction of a stable hybrid star with a pure quark matter core described within two-flavor NJL-type models including color superconductivity, but without strange quarks. This is in strong contradiction to our results where even a phase transition to quark matter seems to be impossible without strange quarks. In Ref.~\\cite{SHH03} this was achieved within a standard NJL model with relatively light up and down quarks, corresponding to a relatively small effective bag constant. The authors of Ref.~\\cite{BGAYT} find stable quark cores if the integrals are regularized by Gaussian form factors, but not for a sharp cut-off. In both papers the hadronic phase is described within relativistic mean field models. These findings call for a systematic survey of the parameter dependence of the results and of the influence of possible extensions and modifications of the model.}.  If there are no stable quark cores in compact stars, color superconductivity is unlikely to exist in any natural surrounding. Nevertheless, the possibility of CS could have an effect on the maximum mass of neutron stars, as is evident from the difference between the solid line and the bold dashed line in Fig.~\\ref{figmass}. While {\\it excluding} CS the maximum mass can be close to a value of two solar masses, {\\it including} CS the maximum mass is reduced to 1.77 solar masses in our calculations. (If we had chosen a smaller value for the quark-quark coupling constant $H$ in \\eq{Lqq} the maximum mass would of course come out somewhere in between these two values.) Introducing hyperons softens the EOS further, and the corresponding maximum mass is further lowered. If the softening is large, as, e.g., in the case of the EOS of Ref.~\\cite{bbs}, the neutron star becomes unstable already before the transition density to a possible quark matter core is reached. In this case the maximum mass falls below the observational limit. \\par In conclusion, the inclusion of CS in the quark phase keeps the neutron star maximum mass well below two solar masses, independently of the details of the hadronic EOS. The observation of a neutron star with a mass well above two solar masses would seriously question our present view on the EOS of quark matter and on the structure of the quark matter phase. \\\\[2mm] {\\bf Acknowledgments:}\\\\ M.B. thanks INFN for financial support during his visit to Catania in November 2002 and Marcello Baldo and Fiorella Burgio for their warm hospitality. M.O. acknowledges support from the Alexander von Humboldt foundation as a Feodor-Lynen fellow."
    },
    "0212/astro-ph0212263_arXiv.txt": {
        "abstract": "CNO abundances in galaxies bear on issues of galactic evolution as well as stellar evolution and nucleosynthesis. Knowledge about them in dwarf and spiral galaxies depends mainly on emission lines from H~{\\sc ii} regions, with information from stars and supernova remnants available for galaxies in the Local Group.  Oxygen abundances can be related to both local and global properties of the parent galaxies % In  H {\\sc ii} regions they range from 1/30th solar (using the new calibration) to slightly above solar. C/O is a more or less smoothly increasing function of O/H, flat at low metallicities and rising above 12+log(O/H) = 8.0 with a 45$^{\\circ}$ slope. N/O behaves quite similarly in the traditional `primary followed by secondary' style, but with a substantial scatter at least around O-abundance 8.0.      C/N ratios show no clear trend with metallicity. ",
        "introduction": "CNO abundances test stellar yields --- C and N from a wide range of stellar masses, and oxygen mainly from massive stars above 10$M_{\\odot}$.  As predicted by Andr\\' e Maeder (1992) and discussed recently by Henry et al.\\ (2000), carbon yields increase with metallicity somewhat like a traditional `secondary' element because of stellar mass loss and nitrogen perhaps even more steeply as a secondary product of the CNO cycle on carbon, although there is also a `primary' component at low metallicities. Oxygen yields decrease with metallicity, at least to some extent, again because of stellar mass loss.  Apart from the implications for stellar evolution and nucleosynthesis, oxygen is the standard element for defining `metallicity' in gas-rich galaxies and so is used also as a test of galactic chemical evolution and as a standard against which to judge abundances of other elements such as iron in stars and C and N in H {\\sc ii} regions.  How well do we know CNO abundances in external galaxies? Most data come from H {\\sc ii} regions and (in the Local Group) supernova remnants. For  H {\\sc ii} regions  of sufficient surface brightness and not too large metallicity, the standard method is to deduce the electron temperature in the O$^{++}$ zone from the  line ratio 4363/5007, use models to find that in the O$^ +$ zone and add the resulting abundances of O$^ +$ and O$^{++}$ relative to H$^ +$. This is a lower limit, because of $\\delta t^ 2$ (Peimbert 1967). Recombination lines are insensitive to temperature (Esteban et al.\\ 2002), but they are weak and the interpretation not completely straightforward (Tsamis et al.\\ 2002). In most cases the effect appears to be small, probably under 0.1 dex, but there is another correction for oxygen locked on grains, which again might approach 0.1 dex.  When 4363 is too weak to measure, as is normally the case when metallicity approaches or exceeds solar, uncertainties multiply. Searle (1971) had the fundamental insight that the spectral changes exhibited by giant H {\\sc ii} regions across Scd galaxies (Aller 1942) were due to a radial abundance gradient, and thus inspired various `empirical' methods of abundance estimation based on strong nebular lines alone. The most widely used is the notorious $R_{23}$ method (Pagel et al.\\ 1979, 1980), calibrated against a model for S5 in M 101 by Shields \\& Searle (1978).  Investigations since then have shown that the oxygen in S5 was overestimated (Kinkel \\& Rosa 1994), as it has been in some other regions (Castellanos et al.\\ 2002), so that our calibrations have overestimated oxygen abundances near solar, and attempted recalibrations by McCall et al.\\ (1985), Evans \\& Dopita (1985) and McGaugh (1991) have various weaknesses of their own, quite apart from the inevitable non-uniqueness of the relationship (cf.\\ Pilyugin 2000). There is some promise in newly developed strong-line indices like S$_{23}$ (D\\'{\\i}az \\& Perez-Montero 2000), S$_{234}$ (Oey \\& Shields 2000)  and especially in refinements of the $R_{23}$ method by Pilyugin (2000, 2001a) in which the relative strengths of [O  {\\sc ii}] and [O {\\sc iii}] lines are taken into account in an empirical manner, reducing the scatter.  Until more of these refinements have been applied, however, abundances near and above solar have to be taken with a grain of salt.  After oxygen, the next step is usually to estimate the N/O ratio. A common assumption, supported by many models at least at low metallicities, is that nitrogen is ionized to the same degree as oxygen (e.g.\\ Garnett 1990), so that N/O = N$^ +$/O$^ +$. This can actually be checked from infra-red data when available, e.g.\\ in 30 Doradus (Lester et al.\\ 1987). Carbon lacks prominent emission lines in the visible or near IR, so one has to use UV observations, usually of the intercombination lines C {\\sc iii}] 1909 and C {\\sc ii}] 2326.  CNO abundances are also now being measured in stars, mostly B stars and supergiants, in the nearest galaxies. As in the Milky Way, comparisons between stellar and nebular abundances provide a useful check. Here one has to be rather careful for two reasons. \\begin{enumerate} \\item Stellar abundances are often expressed differentially relative to the Sun, but the solar goalposts are shifting, at least in the cases of C and O (Holweger 2001; Allende Prieto et al.\\ 2001, 2002) and this has caused some confusion.  \\item In some hot main-sequence stars and supergiants, especially if rotating, the atmospheres have undergone some CNO processing, so to that extent it is the differences between stars and H  {\\sc ii} regions that are of interest; but oxygen is usually unaffected. \\end{enumerate} ",
        "conclusions": ""
    },
    "0212/astro-ph0212055_arXiv.txt": {
        "abstract": "For an accretion disk around a black hole, the strong relativistic effects affect every aspect of the radiation from the disk, including its spectrum, light-curve, and image. This work investigates in detail how the images will be distorted, and what the observer will see from different viewing angles and in different energy bands. ",
        "introduction": "\\label{section:redshiftmaps}  Figures~\\ref{fig:redshiftmap}, \\ref{fig:overallshape} and \\ref{fig:inner_contour} are all redshift maps. Figures~\\ref{fig:redshiftmap} shows the redshift maps of disks with an outer radius $20R_{\\mathrm g}$ ($R_{\\mathrm g} \\equiv \\frac{GM}{c^2}$) and an inner radius as the last stable orbit. For a small viewing angle, there is only significant (gravitational) red-shift in the inner region; whereas for large viewing angles, Doppler shift becomes important, significant blue- and red-shift can be seen on the whole disk. (It would be clearer in color figures.) Also, for large viewing angles, far end of the disk bends toward the observer.  \\vskip -0.1in \\begin{figure}[ht] \\centering \\begin{minipage}{.48\\textwidth} \\centering \\includegraphics[width=\\textwidth]{redshift_map-bw.ps} \\end{minipage} \\begin{minipage}{.5\\textwidth} \\centering \\caption{ Red-shift map and overall shape of accretion disks around black holes. $g$ is the ratio between the observed energy of a photon and the energy in the rest frame of the disk. Top to bottom: $\\theta=5\\deg, 45\\deg, 70\\deg, 85\\deg$, $\\theta$ is the inclination angle (the angle between the normal of the disk and the line of sight). Left to right: $a=0.01, 0.5, 0.998$, where $a$ is the specific angular momentum of the black hole. The white zones stand for the regions with zero red-shift (following Fanton \\etal (1997) ). Left-hand side of the disk is approaching the observer and blue-shifted. } \\end{minipage} \\label{fig:redshiftmap} \\end{figure} \\vskip -0.1in  Figure~\\ref{fig:overallshape} shows the overall shape of disks due to Doppler and gravitational effects at a large inclination angle ($85\\deg$). The inner radius takes the value of the last stable circular orbit. In the first row, the white circles and ellipses show the last stable orbit. The inner contours are greatly distorted by the focusing effect.  Figure~\\ref{fig:inner_contour} shows the comparison between Shwarzschild black holes and Kerr black holes, which shows clearly additional distortion due to frame-dragging effect. Spin results in a smaller last stable orbit, and the inner contour changes from symmetric for $a=0.01$ to asymmetric for $a=0.998$. Therefore in principle, we can tell a rotating black hole from a non-rotating one from the inner contours of their accrection disks.  \\vskip -0.1in \\begin{figure}[ht] \\centering \\begin{minipage}[t]{.48\\textwidth} \\includegraphics[width=\\textwidth]{fig2.eps} \\end{minipage} \\begin{minipage}[t]{.48\\textwidth} \\includegraphics[width=\\textwidth]{fig3.eps} \\end{minipage} \\sidebyside {\\caption{ Redshift maps for disks at inclination angle $85\\deg$ with different outer radii. Top to bottom: disks with outer radii $10R_{\\mathrm g}$, $20R_{\\mathrm g}$, $100R_{\\mathrm g}$. Left to right: $a=0.01, 0.5, 0.998$. See text for details. } \\label{fig:overallshape}} {\\caption{ Redshift maps for 10 $R_{\\mathrm g}$  disks with different inner radii. Top row: inner radii are the radii of the last stable orbits; Bottom row: inner radii are 6 $R_{\\mathrm g}$. Left to right: $a=0.01, 0.5, 0.998$. } \\label{fig:inner_contour}} \\vskip -0.1in \\end{figure} ",
        "conclusions": "\\label{section:summary} We studied the images of standard thin accretion disks around black holes under the strong relativistic effects. The expected images for systems with different black hole spin, viewing from different viewing angles and in different energy bands are given. The spin of the black hole not only determines how close the disk can extend to the center, but also causes additional distortion due to the dragging of the frame.  The secondary image are left out because in most cases it will be blocked by the disk, though it might be important in some cases.  \\noindent{\\bf Acknowledgments}  The authors would like to thank Dr. Fanton for providing us his code of calculating the redshift map and line profile of a thin disk. This study is supported in part by NASA's Marshall Space Flight Center and through NASA's Long Term Space Astrophysics Program.  \\begin{chapthebibliography}{1} \\bibitem{Cadez98} Cadez, A., C. Fanton, and M. Calvani, M. (1998). {\\it New Astronomy}, 3, (No. 8), 647  \\bibitem{Fanton97} Fanton, Claudio, Massimo Calvani, Fernando de Felice, and Andrej Cadez. (1997). {\\it PASJ}, 49, 159  \\bibitem{Luminet78} Luminet, J.-P.. (1978). {\\it A\\&A}, 75, 228  \\bibitem{Page74} Page, D.N. \\& K.S. Thorne. 1974. {\\it ApJ}, 191, 499  \\bibitem{Shakura73} Shakura, N. I. \\& Sunyaev, R. A.. 1973. {\\it A\\&A}, 24, 337  \\end{chapthebibliography}  {"
    },
    "0212/astro-ph0212325_arXiv.txt": {
        "abstract": "We present submillimeter observations of SO, \\soo, \\hhs, \\hhcs, OCS, NS and HCS$^+$ toward nine massive young stars.  The outflow contributes $\\approx$50\\% to the SO and \\soo\\ emission in $15-20''$ beams, more than for CS, where it is 10\\%.  The \\soo\\ abundance increases from dark cloud levels in the outer envelope ($T<100$~K) to levels seen in hot cores and shocks in the inner envelope ($T>100$~K).  Molecular abundances are consistent with a model of ice evaporation in an envelope with gradients in temperature and density for a chemical age of $\\sim 30000$~yr. The high observed abundance of OCS, the fact that \\txc (OCS)$\\gg$\\txc (\\hhs), and the data on solid OCS and \\hhs\\ all suggest that the major sulphur carrier in grain mantles is OCS rather than \\hhs.  For most other sulphur-bearing molecules, the source-to-source abundance variations by factors of up to 10 do not correlate with previously established evolutionary trends in temperature tracers.  These species probe the chemically inactive outer envelope.  Our data set does not constrain the abundances of \\hhs\\ and SO in the inner envelope, which, together with \\soo, are required to use sulphur as a clock. ",
        "introduction": "\\label{sec:intro}  Spectral line surveys at submillimeter wavelengths have revealed considerable chemical differences between star-forming regions (see \\cite{evd01} for an overview). While these differences indicate activity, the dependence of molecular abundances on evolutionary state and physical parameters is poorly understood. Better insight into this relation would be valuable for probing the earliest, deeply embedded phases of star formation where diagnostics at optical and near-infrared wavelengths are unavailable. This is especially true for the formation of high-mass stars, for which the order in which phenomena occur is much less well understood than in the low-mass case, and for which the embedded phase is a significant fraction ($\\approx$10\\%) of the total lifetime of $\\sim$10$^6$~yr.  \\begin{figure} \\begin{center} \\resizebox{\\hsize}{!}{\\includegraphics[height=10cm,angle=-90]{vandertakf2_1.ps}} \\caption{Examples of the spectra taken with the JCMT. The top and bottom frequency scales of the left panel are the two receiver sidebands; the right panel only shows lines from one sideband. Adopted from van der Tak et al.~(2002).} \\label{fig:data} \\end{center} \\end{figure}  Sulphur-bearing molecules are attractive as candidate tracers of early protostellar evolution (\\cite{jh98sulf}). Models of ``hot core'' chemistry (\\cite{char97}), as well as of shock chemistry (\\cite{gpdf93}) predict strong variations in the abundances of sulphur molecules on time scales of $\\sim$10$^4$~yr. To test these scenarios, this paper discusses submillimeter observations of sulphur-bearing molecules towards nine regions at the earliest stages of high-mass star formation, with luminosities $1\\times 10^4 - 2\\times 10^5$~\\lsol\\ at distances of $1-4$~kpc. This paper only summarizes our results; a complete description is given by \\cite*{fvdt02}. ",
        "conclusions": ""
    },
    "0212/astro-ph0212113_arXiv.txt": {
        "abstract": "{ In a companion paper, we investigated the question of the spatial origin of the cosmic rays detected in the Solar neighborhood, in the case of standard sources located in the Galactic disk. There are some reasons to believe that there may also be a large number of sources located in the halo, for example if the Galactic dark matter is made of supersymmetric particles or if Primordial Black Holes are present. These exotic sources could enhance the $\\bar{p}$, $\\bar{d}$ or positrons above the standard background, indicating the existence of new physics. The spatial distribution of these hypothetical sources, though an important ingredient to evaluate these exotic signals, is poorly known. The aim of this paper is to point out that this discussion should not be disconnected from that of the propagation properties in the Galaxy. More precisely, we determine the regions of the halo from which a significant fraction $f$ of cosmic rays antiprotons and antideuterons detected in the Solar neighborhood were emitted (we refer to these regions as $f$-volumes), for different sets of propagation parameters consistent with B/C data, as derived in \\citet{Maurin02b}. It is found that some of them lead to rather small $f$-volumes, indicating that the exotic cosmic rays could have a local origin (in particular for a small halo or a large Galactic convective wind), coming from the solar neighborhood or the Galactic center region. It is also found that the dark matter density enhancement (spike) due to the accretion around the central supermassive black hole gives a negligible contribution to the exotic charged particle signal on Earth. The case of electrons and positrons is also discussed. ",
        "introduction": "A great amount of work has been done these last twenty years on the astrophysical signatures that could unravel new physics. In the eighties, there were great hopes that the antiproton signal, which showed an excess at an energy of a few hundreds of MeV in the first balloon experiments, could be such a signature. However, this hope was swept away by the progress in measurements -- see e.g. {\\sc bess} \\citep{Orito00,Maeno01} or {\\sc heat} \\citep{Beach01} and {\\sc caprice} \\citep{Boezio01} at higher energy -- and a better determination of the cosmic ray propagation parameters (see e.g. \\citet{Maurin02b}). It was shown that the measured antiproton flux was indeed compatible with the sole secondary standard spallative production \\citep{Bergstrom99b,Donato01}  (see the first paper for a comprehensive historical discussion and a panel of references dealing with exotic antiproton production).  \\citet{Donato00} showed that the antideuteron ($\\bar{d}\\,$) signal could lead to a clearer signature of SUSY. However, as discussed in many other studies on SUSY antiprotons \\citep{Rudaz88,Stecker89,Jungman94,Bottino95,Bottino98,Wells99,Bergstrom99b}, the indeterminacy in the dark matter distribution, as well as its possible clumpiness \\citep{Bergstrom99}, might severely change the conclusions. In contrast, the Hawking evaporation of Primordial Black Holes (PBH) could also yield a new source of cosmic rays \\citep{Maki96}, but the precise shape of the dark matter in this case is not crucial \\citep{Barrau02,Barrau03}. Nevertheless, in the latter case, it was shown that even considering only the propagation parameters giving a good fit to B/C data, the remaining degeneracy for example in the diffusive halo height has sizeable effects on the primary flux \\citep{Barrau02}.  Hence, at least two different phenomena can affect the conclusions reached in papers dealing with exotic flux calculations. The first one, related to the spatial distribution of SUSY sources, is usually thoroughly discussed \\citep{Bergstrom99b}, but the second point - namely the influence of various propagation parameters - is generally skipped, due to the simplicity of the propagation models used. The aim of the paper is {\\em not} to compare the predicted $\\bar{p}$, $\\bar{D}$ fluxes with observations for different series of models, but rather to point out which characteristics of the models actually play a role, in order to give some physical insights and milestones for studies specifically devoted to exotic flux evaluations.  We apply the method described in \\citet{Taillet03} to determine the volumes from which a fraction $f$ of cosmic rays reaching the Solar neighborhood were emitted, or equivalently the volumes that contribute to the fraction $f$ of the total flux detected in the Solar neighborhood. These volumes will be referred to as the $f$-volumes throughout the paper.  We find that depending on the diffusion parameters (evaluated from a systematic study of standard CR, \\citet{Maurin02b}) as well as on the source spatial distribution, the spatial origin of cosmic rays may be quite local, the particles detected in the Solar neighborhood having mostly been created a few kpc away from the Solar neighborhood in some cases, or a few kpc away from the Galactic center in others. ",
        "conclusions": "This paper analyzes the spatial origin of exotic particles created from a dark matter profile. We presented the $f$-volumes inside which a given fraction of the cosmic rays detected in the Solar neighborhood were emitted. At high energy ($E \\gg 1$ GeV/nuc), the shape of the isodensity surfaces  is set by the geometry of the diffusive halo, in particular on its height $L$, the influence of the side boundary at $r=R$ being small. We then showed that the $f$-volumes defined are smaller when spallations and convection are taken into account, but in a very different way: for particles in the diffusive halo, the wind exponentially decreases the probability of reaching the Galactic plane, whereas spallations have about a null effect on the latter. The parameters $L$ and $2V_c/K$ indicate whether the propagation is convection or escape-dominated. In Table~\\ref{tab5} we summarize the parameters that act as a cut-off in various situations. \\begin{table*}[ht] \\begin{center} \\begin{tabular}{|c||cc|c|}   \\hline Cut-off& Escape-dominated & Convection-dominated & Losses-dominated\\\\ & ($\\chi_{\\rm w}\\gg1$)\t& ($\\chi_{\\rm w}\\ll1$) & ($e^-e^+\\gtrsim$~GeV)\\\\\\hline\\hline Halo& $L$\t\t& $L^*\\approx 3K/V=3r_{\\rm w}/2$ & $ \\approx 5r_{\\rm loss}$\\\\ Radial& $\\min(R,3L)$ & $\\min(R,3L^*)$ &  $\\approx 5 r_{\\rm loss}$\\\\\\hline \\end{tabular} \\caption{Summary of the typical cut-off in $z$ and $r$ directions beyond where a cosmic ray cannot originate. The sole parameters that determine these cut-offs are $L$ (halo size) and/or $\\chi_{\\rm w}\\equiv r_{\\rm w}/L$ -- related to the convective wind $V_c$ {\\em via\\/} $r_{\\rm w}=2K/V_c$, or $r_{\\rm loss}$ related to the effective life-time for positrons and electrons.} \\label{tab5} \\end{center} \\end{table*} Two source distribution for the isothermal dark matter profile were considered: production related to the density of the source (e.g. PBH evaporation), or production related to the square of the density of the sources (e.g. SUSY annihilation). The 99\\%-volumes are strongly stretched toward the Galactic center, corresponding to the maximum of the source distribution. This follows from the competition between the effective source which is maximum at the Galactic center, and the probability density which steadily decreases from our position $R_{\\odot}$ to reach $\\sim 10^{-4}-10^{-5}$~kpc$^{-3}$ for purely diffusive regime (or even less when convection is included) at the Galactic center. The fluxes in the Solar neighborhood are found to be far more sensitive to the dark matter profile in the SUSY case than in the PBH case. In both cases, the side boundary of the diffusive volume is observed to play a negligible role as long as $R\\gtrsim 20-30$~kpc.  As a last step, realistic propagation parameters were implemented, and the key parameters were found to be the halo size $L$ and the diffusion slope $\\delta$ (actually $V_c/K_0$). For the species considered here (antiprotons and antideuterons), spallations always play a negligible role in the origin. It was found that this origin is far more local in case of large $\\delta$ and small $L$ than in case of small $\\delta$ and large $L$. Moreover, the shape of the dark matter distribution near the Galactic center does not matter so much for the PBH case, whereas it may be crucial for SUSY annihilating particles. We emphasized that in any discussion of the annihilation signal in charged particles, the propagation parameter $\\delta$ or more precisely, the presence of a Galactic wind, should be considered, with the same importance of the parameter $L$ or the choice of the dark matter profile.  Two last points are worth noting. First, even though the work presented here does not allow a quantitative estimation of the effect of possible clumpiness of the dark matter halo (for SUSY annihilations), we observed that the comparison between the electron and antiproton SUSY signals should involve a careful inspection of the corresponding boost factors. Second, whereas the use of B/C-induced propagation parameters is justified for standard antiprotons (corresponding $f$-surfaces can be seen in \\citet{Taillet03}), there is no guarantee that these parameters are valid in the $f$-volumes depicted here."
    },
    "0212/astro-ph0212396_arXiv.txt": {
        "abstract": "The detailed morphology of the interstellar medium (ISM) in the central kiloparsec of galaxies is controlled by pressure and gravitation. The combination of these forces shapes both circumnuclear star formation and the growth of the central, supermassive black hole. We present visible and near-infrared {\\it Hubble Space Telescope} images and color maps of 123 nearby galaxies that show the distribution of the cold ISM, as traced by dust, with excellent spatial resolution. These observations reveal that nuclear dust spirals are found in the majority of active and inactive galaxies and they possess a wide range in coherence, symmetry, and pitch angle. We have used this large sample to develop a classification system for circumnuclear dust structures. In spite of the heterogeneous nature of the complete sample, we only find symmetric, two-arm nuclear dust spirals in galaxies with large scale bars and these dust lanes clearly connect to dust lanes along the leading edges of the large scale bars. Not all dust lanes along large scale bars form two arm spirals, however, and several instead end in nuclear rings. We find that tightly wound, or low pitch angle, nuclear dust spirals are more common in unbarred galaxies than barred galaxies. Finally, the extended narrow line region in several of the active galaxies is well-resolved. The connection between the ionized gas and circumnuclear dust lanes in four of these galaxies provides additional evidence that a significant fraction of their extended narrow line region is ambient gas photoionized {\\it in situ} by the active nucleus. In a future paper, we will use our classification system for circumnuclear dust to identify differences between active and inactive galaxies, as well as barred and unbarred galaxies, in well-matched subsamples of these data. ",
        "introduction": "The kinematics of the cold interstellar medium (ISM) largely determines the future star formation and black hole growth in galaxies. In the circumnuclear region of most galaxies, corresponding to approximately the central kiloparsec, both pressure and gravitational forces can dominate these kinematics. In the absence of velocity information, observations of dust lanes in absorption from visible and near-infrared (NIR) images can provide a snapshot of the cold ISM.  Before the advent of \\hst, high spatial resolution observations of dust in the circumnuclear region were only possible for relatively nearby galaxies, including our own \\citep{morris96}, Andromeda \\citep{mcelroy83,ciardullo88}, and several other relatively nearby galaxies where circumnuclear dust is apparent on photographic plates in, for example, the Carnegie Atlas of Galaxies \\citep{sandage94}. Large single-band \\hst programs have vastly expanded the number of galaxies with high spatial resolution observations \\citep[resolutions less than 50 parsecs;][]{malkan98,carollo98,ferruit00}. These observations have revealed a wealth of detail in the circumnuclear structure of spiral galaxies and show that this nuclear structure is frequently quite different from the large scale morphology in any given galaxy. Multiband observations of smaller samples of active galaxies clearly show that much of this circumnuclear structure is due to nuclear dust \\citep{elmegreen98,quillen99}, although mostly material that is only a factor of a few higher density than the ambient ISM \\citep{regan99,martini99}.  One of the most common dust features found in the circumnuclear region are nuclear dust spirals, which exhibit a wide range of intrinsic shapes. These include symmetric, two-arm spirals, one-arm spirals, and chaotic, multiarm spirals \\citep{elmegreen98,laine99,regan99,martini99}. A number of theoretical models exist that can give rise to coherent dust structures in the circumnuclear region, although in all cases some differential rotation is required to form the spiral structure. Galaxies with strong, large scale bars form shock fronts along the leading edges of their bar \\citep{prendergast83,athanassoula92} and hydrodynamic simulations show that this inflowing material can form symmetric and two-arm or ``grand design'' spiral structure in the circumnuclear region \\citep{englmaier00,patsis00}. \\citet{maciejewski02} recently modeled the circumnuclear region of barred galaxies and found grand design nuclear spirals similar to those found by \\citep{englmaier00}. While the nuclear spirals in \\citet{englmaier00} only corresponded to weak perturbations, the gas velocity field of the spirals in \\citet{maciejewski02} had a large negative divergence, characteristic of strong shocks and similar to the shocks in the large-scale bar. Significant inflow is likely to be associated with these spirals shocks, as shown by \\citet{fukuda98}.  More chaotic, or ``flocculent'' spirals can form in the circumnuclear region from pressure forces, or acoustic turbulence, in the ISM \\citep{elmegreen98, montenegro99}. Detailed studies of the LINERs NGC\\,4450 and NGC\\,4736 by \\citet{elmegreen02} show that the azimuthal Fourier transform spectra have a power law slope of -5/3, characteristic of turbulence. These structures therefore appear to be formed by the same turbulent processes that create structure in the ISM of galaxies at larger radii, but have been sheared by the presence of some differential rotation in the circumnuclear region and given a spiral appearance. Recent hydrodynamic simulations of the central ISM also naturally produce a filamentary, spiral structure in surface density with pressure and gravitational forces \\citet{wada99,wada01}. Both these analytic and numerical studies suggest some inflow can occur due to turbulence alone.  The nuclear dust spirals on these small scales (they are generally a kiloparsec or less in length) are distinct from their host galaxy spiral arms both spatially and in composition. Nuclear spirals rarely extend out of the region dominated by the central bulge. Observations of Seyfert galaxies over many kiloparsecs show that while the grand design nuclear spirals appear to connect to the dust lanes along the large scale bar, most types of nuclear spirals do not connect to the large-scale spiral arms \\citep{pogge02}. These nuclear spirals also appear to only be enhancements in the dust surface density and not sufficiently overdense to be self-gravitating. Only a small number of nuclear dust spirals have associated star formation \\citep[\\eg NGC\\,4321, Coma D15:][]{knapen95,caldwell99} and NIR observations do not show enhancements in the stellar surface density associated with the spirals \\citep{martini01}. In addition, density waves tend to get stronger with increasing radius, whereas the nuclear spirals get weaker with increasing radius \\citep{montenegro99,elmegreen02}.  One of the primary motivations for most of these multiband observations was to study the circumnuclear environment of active galaxies in order to determine the processes that removed angular momentum from host galaxy material and fueled the central, supermassive black hole. Recent \\hst programs have provided strong evidence that all galaxies with a spheroid component have a supermassive black hole \\citep{richstone98,ferrarese00, gebhardt00,ferrarese01,gebhardt01}. This result has added new urgency to the question of how nuclear activity is fueled and why it is not currently fueled in all galaxies. Nuclear spiral structure was proposed as a mechanism for fueling nuclear activity because it was both discovered in nearly all of the observations of active galaxies and that these structures could be formed by shocks \\citep{regan99,martini99}.  Both bar-driven inflow and acoustic turbulence are theoretically viable mechanisms for the removal of angular momentum. In the case of inflow from large-scale bars, the gas and dust loses essentially all of its angular momentum in the shock front at the leading edge of the bar. For acoustic turbulence, the shocks in the ISM caused by the turbulence could dissipate sufficient energy and angular momentum to lead to significant inflow \\citep{elmegreen02}. However, aside from the absence of kinematic data to demonstrate that inflow does occur, the selective investigation of active galaxies resulted in relatively few comparable observations of inactive galaxies. Such observations are needed to determine if there is an excess of nuclear spiral structure, or other morphological features, in active galaxies over inactive galaxies of similar Hubble types and luminosities. We have obtained \\hst observations of a significantly larger sample of both active and inactive galaxies to perform this analysis and present the data and a morphological study of each of the galaxies in our sample. In a future paper \\citep{martini02} we compare the circumnuclear environment of active and inactive galaxies for a well-matched subsample of these data, as well as discuss some interesting results on differences in the circumnuclear environments of barred and unbarred galaxies. ",
        "conclusions": "We have presented \\hst observations of circumnuclear dust structure in a large sample of nearby galaxies. The color maps which trace this dust structure show that the nuclear dust spirals reported previously are the most common morphological feature in the circumnuclear region of spiral galaxies, irrespective of the presence or absence of an active nucleus or a large-scale bar. We have defined a set of four different classes of nuclear spiral structure that together encompass the main morphological features of nuclear spiral structure. There are in addition many galaxies that show no evidence of nuclear spiral structures and these galaxies have either very chaotic circumnuclear dust or no clearly discernible dust structure in their circumnuclear environment.  While the entire sample is too heterogeneous for a detailed examination of the frequency of nuclear spiral structure in active and inactive galaxies beyond the statement that it is a common feature of both types, there are several results, including two clear differences between barred and unbarred galaxies, that should not be affected by the various selection effects in this sample:  \\noindent 1. Only galaxies with large scale bars exhibit grand design nuclear spiral structure. Where high-quality data on larger scales is available, the two arms of the grand design nuclear spiral connect to the dust lanes on the leadings edges of the large scale bar. As these cases demonstrate that some large scale bars can remove sufficient angular momentum to transport matter to within at least tens of parsecs of the nucleus, they provide qualitative support to the hypothesis that some large scale bars can directly fuel nuclear activity. However, these dust lanes are also found in galaxies that are not known to harbor AGN.  \\noindent 2. Tightly wound nuclear spiral arms are more likely to be found in unbarred galaxies than barred galaxies.  \\noindent 3. Several active galaxies with prominent ionization cones, most notably Mrk\\,573, Mrk\\,1066, and NGC\\,788, show that the emission line gas is the same material that forms the spiral dust lanes in the circumnuclear region. These dust lanes, traced by their absorption of stellar light, become bright emission regions where they cross the opening angle of the ionization cone from the nucleus.  In a future paper we will compare a subsample of active galaxies to an inactive, control sample matched in Hubble type, luminosity, heliocentric velocity, size, and inclination to determine if there are any differences in the circumnuclear dust structures between active and inactive galaxies. We will also employ a similar sample-matching strategy to verify these differences between the circumnuclear dust structures of barred and unbarred galaxies to extend the results described above."
    },
    "0212/astro-ph0212443_arXiv.txt": {
        "abstract": "We demonstrate a close analogy between a viscoelastic medium and an electrically conducting fluid containing a magnetic field. Specifically, the dynamics of the Oldroyd-B fluid in the limit of large Deborah number corresponds to that of a magnetohydrodynamic (MHD) fluid in the limit of large magnetic Reynolds number.  As a definite example of this analogy, we compare the stability properties of differentially rotating viscoelastic and MHD flows. We show that there is an instability of the Oldroyd-B fluid that is physically distinct from both the inertial and elastic instabilities described previously in the literature, but is directly equivalent to the magnetorotational instability in MHD.  It occurs even when the specific angular momentum increases outwards, provided that the angular velocity decreases outwards; it derives from the kinetic energy of the shear flow and does not depend on the curvature of the streamlines.  However, we argue that the elastic instability of viscoelastic Couette flow has no direct equivalent in MHD. ",
        "introduction": "\\subsection{Viscoelastic and magnetohydrodynamic fluids}  In his investigation of the viscosity of gases, Clerk Maxwell (1867) proposed that the stress in a fluid obeys an equation of the form \\begin{equation} (\\mbox{stress})+\\tau{{{\\rm d}(\\mbox{stress})}\\over{{\\rm d}t}}= (\\mbox{viscosity})\\times{{{\\rm d}(\\mbox{strain})}\\over{{\\rm d}t}}, \\label{maxwell} \\end{equation} where $\\tau$ is the relaxation time.  If the time scale of the straining motion is long compared to $\\tau$, the second term on the left-hand side is negligible and the stress is proportional to the rate of strain.  This Newtonian relation gives rise to viscous behaviour.  However, if the time scale of the strain is short compared to $\\tau$, the first term on the left-hand side is negligible and the stress is proportional to the strain itself.  This Hookean relation gives rise to elastic behaviour.  The reason for the elastic response is that the rapid strain prevents the configuration of the molecules from relaxing towards an equilibrium distribution, and the stress is therefore `frozen in' to the fluid.  Modern constitutive equations for viscoelastic fluids (Bird, Armstrong \\& Hassager 1987{\\it a\\/}) are usually expressed in a covariant tensorial form based on the principles set out by Oldroyd (1950).  His liquid B is one of the most widely used nonlinear models of a viscoelastic fluid, and provides a fair representation of a dilute solution of a polymer of high molecular weight.  It is based on Maxwell's equation (\\ref{maxwell}), but cast in a form that satisfies the principle of material frame indifference.  Moreover, it can also be derived from the kinetic theory of idealized extensible polymer molecules contained in a Newtonian solvent (Bird, Curtiss \\& Armstrong 1987{\\it b\\/}).  The dimensionless number characterizing the ratio of the relaxation time to the time scale of the flow is the Deborah number (or Weissenberg number); when this is large, the polymeric stress is effectively `frozen in' to the fluid.  In an electrically conducting fluid, the magnetic field $\\bB$ affects the dynamics through the bulk Lorentz force (e.g. Roberts 1967).  This can be represented in terms of the Maxwell electromagnetic stress tensor,\\footnote{Here $\\mu_0$ is the permeability of free space.  For non-relativistic flows, the electric field makes a negligible contribution to the stress.} \\begin{equation} \\maxwell={{\\bB\\bB}\\over{\\mu_0}}-{{B^2}\\over{2\\mu_0}}\\unittensor, \\label{maxwell_stress} \\end{equation} the two parts of which correspond to a tension in the field lines and an isotropic magnetic pressure.  It is well known that, in a perfectly conducting fluid, the magnetic field is `frozen in' to the fluid, in the sense that magnetic field lines can be identified with material lines (Alfv\\'en 1950).  Even in a fluid of finite conductivity, the magnetic field is effectively `frozen in' for motions of sufficiently short time scale, or sufficiently large length scale, corresponding to a large magnetic Reynolds number.  It follows that the Maxwell stress is also `frozen in' to the fluid in a certain sense.  From a mathematical point of view, the `freezing in' of a tensor field $\\tensorX(\\br,t)$ in a flow with velocity field $\\bu(\\br,t)$ can be expressed by the equation (e.g. Tur \\& Yanovsky 1993) \\begin{equation} {{\\partial\\tensorX}\\over{\\partial t}}+\\pounds_{\\bu}\\tensorX=\\zerotensor, \\end{equation} where $\\pounds$ is the Lie derivative.  For a scalar field $X$, this gives the familiar expression $\\partial X/\\partial t+\\bu\\bcdot\\bnabla X=0$, meaning that the numerical value of $X$ is conserved by every fluid element.  For a (contravariant) vector field $\\bB$ it gives the induction equation of ideal, incompressible magnetohydrodynamics (MHD), which implies that the magnetic field is advected and stretched in the same way as infinitesimal line elements.  For a second-rank tensor field it results in the `upper-convected derivative' that appears in the governing equation of the Oldroyd-B fluid (Bird {\\it et al.}  1987{\\it a\\/}).  At a physical level, an analogy is to be seen between a polymer solution, containing extensible molecules that are advected and distorted by the flow and react on it through their tension, and an electrically conducting fluid, containing magnetic field lines that are also advected and distorted by the flow and react on it through their tension.  It follows that there is a physical and mathematical similarity between the dynamics of viscoelastic and MHD fluids.  We will show that a formal analogy can be drawn between the Oldroyd-B fluid in the limit of large Deborah number and an MHD fluid in the limit of large magnetic Reynolds number.  In other words, in this limit, the Maxwell stress in MHD obeys the equation of a Maxwell fluid.  \\subsection{Instabilities of differentially rotating fluids}  \\label{instabilities}  Differentially rotating flows are common in astrophysics and geophysics, and have been studied extensively in the laboratory.  The simplest form of differential rotation occurs when the angular velocity depends only on the cylindrical radius, $\\Omega=\\Omega(r)$, and Couette flow between differentially rotating cylinders provides an excellent model system for investigating the dynamics of such flows. According to Rayleigh (1916), instability occurs in the absence of viscosity whenever the specific angular momentum $|r^2\\Omega|$ decreases outwards.  Most subsequent theoretical and experimental studies, starting with the classic work of Taylor (1923), have focused on the onset of Rayleigh's inertial instability in the presence of viscosity, and the interesting sequence of dynamical states that ensues.  Numerous variants of Couette flow have also been considered, and among these the Couette flow of viscoelastic fluids has received much attention.  Early work in the 1960s (e.g. Thomas \\& Walters 1964; Giesekus 1966) examined the effect of viscoelasticity on the onset of Rayleigh's inertial instability.  However, one of the most important recent results is the theoretical and experimental demonstration by Larson, Shaqfeh \\& Muller (1990) of a physically distinct instability in viscoelastic Couette flow.  This is a purely elastic instability that occurs at sufficiently large Deborah number $\\tau|{\\rm d}\\Omega/{\\rm d}\\ln r|$, even in the limit of negligible inertia, and irrespective of the sign of the angular momentum gradient.  The influence of a magnetic field on the stability of the Couette flow of an electrically conducting fluid has also been investigated theoretically and (to a lesser extent) experimentally.  Again, early work (described by Chandrasekhar 1961) focused on the effect of the magnetic field on the onset of inertial instability.  More importantly, Velikhov (1959) and Chandrasekhar (1960) uncovered a physically distinct instability in magnetized Couette flow.  In the absence of viscosity and resistivity, and in the presence of a weak vertical magnetic field, this `magnetorotational' instability occurs whenever the angular velocity $|\\Omega|$ decreases outwards, irrespective of the sign of the angular momentum gradient.  The magnetorotational instability finds its most important applications in astrophysical fluid dynamics, where magnetic fields are prevalent and the astronomical length scales allow for large magnetic Reynolds numbers.  The stability of differentially rotating flows is of considerable interest in astrophysics, especially in connection with accretion discs (e.g. Pringle 1981).  These are usually thin discs of gas in circular orbital motion around a star or black hole.  The angular velocity decreases outwards according to Kepler's third law, $\\Omega\\propto r^{-3/2}$, and the Reynolds numbers are extremely high (e.g. $10^{14}$).  Observations indicate that angular momentum is transported outwards through accretion discs at a much greater rate than allowed by viscosity, and understanding the origin of this `anomalous viscosity' has been a major goal of accretion disc research.  The anomalous viscosity is usually attributed to turbulent transport. However, despite the very high Reynolds numbers, there is no convincing demonstration of any suitable hydrodynamic instability in circular Keplerian flow.  Indeed, simple reasoning can be used to argue that hydrodynamic turbulence is unlikely to be self-sustaining in a flow that amply satisfies Rayleigh's stability criterion (Balbus \\& Hawley 1998).  However, it has been demonstrated that MHD turbulence develops very readily in accretion discs, through the nonlinear development of the magnetorotational instability.  Since the results of Velikhov (1959) and Chandrasekhar (1960) were rediscovered by Balbus \\& Hawley (1991) and their significance was appreciated, the magnetorotational instability has been analysed in considerable detail in the astrophysical literature.  \\subsection{Properties of the magnetorotational instability}  \\label{mri}  The properties of the magnetorotational instability have been reviewed by Balbus \\& Hawley (1998), and we recall some of the important features here.  Its simplest manifestation is in an incompressible, inviscid, perfectly conducting fluid having angular velocity $\\Omega(r)$ and containing a uniform magnetic field $\\bB$ parallel to the axis of rotation.  An approximate local dispersion relation can be obtained for normal modes having growth rate $s$ and wavevector $\\bk$ parallel to $\\bB$, and has the form \\begin{equation} s^4+s^2\\left[4\\Omega(\\Omega-A)+2\\omega_{\\rm A}^2\\right]+ \\omega_{\\rm A}^2(\\omega_{\\rm A}^2-4\\Omega A)=0, \\label{magnetorotational_dispersion_relation} \\end{equation} where the quantity \\begin{equation} A=-{{r}\\over{2}}{{{\\rm d}\\Omega}\\over{{\\rm d}r}} \\end{equation} measures the differential rotation, and is known as Oort's first constant in the astrophysical literature.  The Alfv\\'en frequency is $\\omega_{\\rm A}=(\\mu_0\\rho)^{-1/2}\\bk\\bcdot\\bB$, and the combination \\begin{equation} 4\\Omega(\\Omega-A)={{1}\\over{r^3}}{{\\rm d}\\over{{\\rm d}r}}(r^4\\Omega^2) \\end{equation} is the Rayleigh discriminant, or the square of the epicyclic oscillation frequency.  When this is positive, unstable normal modes with $s>0$ can nevertheless be found provided that $4\\Omega A>0$, as is the case in astrophysical discs.  The maximal growth rate, $s=|A|$, is achieved by a mode having $\\omega_{\\rm A}^2=A(2\\Omega-A)$.  More generally, and from a local perspective, rotation and shear in the correct relative orientation are required, and a weak magnetic field of any geometry is sufficient to initiate the instability, provided the fluid is sufficiently ionized.  In the presence of significant dissipation, the ideal growth rate must compete with viscous and resistive damping, so that growth rates less than $|A|$ are achieved, or the instability may be suppressed altogether.  An unstable mode must always bend the field lines, having a non-zero Alfv\\'en frequency, and therefore the instability of a purely azimuthal (or toroidal) field is essentially non-axisymmetric.  A simple explanation of the instability can be given in terms of two fluid elements, connected by magnetic field lines, that are initially in circular orbit at the same radius.  The fluid elements are then given angular momentum perturbations of opposite sign.  The one receiving the positive perturbation moves to an orbit of larger radius and acquires a smaller angular velocity, lagging behind its partner. The tension of the magnetic field exerts a torque that pulls the lagging element forwards, enhancing the initial perturbation and leading to instability.  A mechanical analogue, consisting of two orbiting particles connected by a weak spring, also exhibits instability.  \\subsection{Plan of the paper}  The main purpose of this paper is to draw attention to the physical and mathematical similarity between viscoelasticity and MHD.  As an example of the application of this idea, we explore in some detail the relation between the instabilities of differentially rotating viscoelastic and MHD flows.  As described in \\S~\\ref{instabilities}, previous investigations have uncovered instabilities of viscoelastic and MHD Couette flow that are physically distinct from Rayleigh's inertial instability.  In the light of the analogy we describe, an obvious question is whether the elastic instability of Larson {\\it et al.} (1990) is somehow related to the magnetorotational instability. We will argue that this is not the case, but will show that there is another instability of viscoelastic Couette flow that is the direct equivalent of the magnetorotational instability.  The remainder of this paper is organized as follows.  In \\S~2 we set out the basic equations governing incompressible viscoelastic and MHD flows and present the analogy between them.  We then discuss, in \\S~3, the possible sources of instability on the basic of energy considerations.  In \\S~4 we define a model system consisting of plane Couette flow in a rotating channel, equivalent to cylindrical Couette flow in the narrow-gap limit, and formulate eigenvalue problems for the normal modes of the Oldroyd-B and MHD fluids.  Some numerical solutions are presented in \\S~5 to illustrate the expected similarity between the two systems.  In \\S~6 the existence of localized growing solutions in the two systems, satisfying the same magnetorotational dispersion relation, is demonstrated by an asymptotic analysis. Finally, the results are summarized and discussed in \\S~7. ",
        "conclusions": "We have demonstrated a close analogy between a viscoelastic medium and an electrically conducting fluid containing a magnetic field.  Both an Oldroyd-B fluid, in the limit of large Deborah number, and a magnetohydrodynamic fluid, in the limit of large magnetic Reynolds number, feature a stress tensor that is nearly `frozen in' to the fluid in a precise mathematical sense.  As a definite example of this analogy, we have examined a local model of a differentially rotating fluid, consisting of plane Couette flow in a rotating channel.  The stress tensor in the case of a viscoelastic fluid resembles the Maxwell stress corresponding to a magnetic field aligned with the flow.  Our analysis demonstrates that there is a detailed correspondence between instabilities in the two systems.  We have identified a direct equivalent of the magnetorotational instability in the viscoelastic fluid.  It exists when the angular velocity and relative vorticity are antiparallel (or when the angular velocity decreases outwards) and the maximal growth rate is equal to the shear parameter, or Oort constant, $A$.  It is distinguished most clearly from Rayleigh's inertial instability by the fact that it occurs even when the specific angular momentum increases outwards.  It is also distinct from the elastic instability described by Larson {\\it et al.} (1990), which depends on the curved geometry of Couette flow and exists in the elastic limit, ${\\it De}/{\\it Re}\\to\\infty$.  We have also found an axisymmetric viscoelastic instability that can be understood as being analogous to the magnetorotational instability of a vertical magnetic field.  This reflects the fact that the polymeric stress tensor $\\stress_{\\rm p}$ can be decomposed into three effective magnetic fields, one of which is a uniform field almost aligned with the flow and another of which is a uniform vertical field.  Although the vertical field is much weaker in the limit of large ${\\it De}$, it provides the dominant restoring force for axisymmetric disturbances.  The instability discussed by Larson {\\it et al.} (1990) does not appear in our analysis because we have neglected the curvature of the streamlines.  Although Larson {\\it et al.} (1990) considered the limit of a narrow gap, the characteristic growth rates they obtained are smaller than the growth rates we have discussed, by a factor of order $\\epsilon^{1/2}$, where $\\epsilon$ is the ratio of the gap width to the radius.  Being purely elastic in nature, the instability of Larson {\\it et al.} (1990) must derive its energy from the elastic energy stored in the flow, rather than the shear energy which is the source for inertial and magnetorotational instabilities.  It is natural to enquire whether there is an MHD equivalent of the instability of Larson {\\it et al.} (1990), in which energy is derived from the magnetic field.  In cylindrical Couette flow at large ${\\it De}$ the dominant stress component has the form $T_{\\phi\\phi}\\propto r^{-4}$, which can be identified with an azimuthal (or toroidal) magnetic field $B_\\phi\\propto r^{-2}$.  Typically, toroidal pinch configurations are unstable to modes that rely on the curvature of the magnetic field lines and derive energy from the magnetic configuration.  The most dangerous are the $m=1$ `kink' modes and $m=0$ `sausage' modes.  However, in the absence of fluid motion the profile $B_\\phi\\propto r^{-2}$ is sufficiently steep to be stable to all perturbations (Tayler 1973).  If the profile of $T_{\\phi\\phi}$ happened to be less steep, for example $T_{\\phi\\phi}\\propto r^{-1}$, it is likely that there would be a viscoelastic equivalent of the kink instability.  We therefore conclude that the instability of Larson {\\it et al.} (1990) is not directly related to an MHD instability, but relies on inherently viscoelastic effects not captured by our analogy. This conclusion is supported by an examination of the physical explanation that Larson {\\it et al.} give for their instability.  Since the work of Larson {\\it et al.} (1990) there have been a number of related theoretical and experimental studies of viscoelastic Couette flow, some of which have examined the effects of inertia and non-axisymmetry (e.g. Avgousti \\& Beris 1993; Steinberg \\& Groisman 1998; Baumert \\& Muller 1999).  Some of these might have been expected to reveal the analogue of the magnetorotational instability.  However, it appears that the cases usually investigated are those in which either the outer or inner cylinder is stationary, or the inner cylinder rotates at twice the angular velocity of the outer cylinder with only a narrow gap between the two.  This means that, whenever the analogue of the magnetorotational instability might have occurred, the system is unstable to Rayleigh's inertial instability.  In order to separate the two effects, it would be valuable to examine cases in which the angular velocity decreases outwards but the specific angular momentum increases.  Interestingly, the magnetorotational instability has never been demonstrated in laboratory experiments.  Although there is currently much effort towards this goal (e.g. Goodman \\& Ji 2002), the technical requirements are considerable and the system is constrained by the very small magnetic Prandtl numbers of liquid metals.  A viscoelastic magnetorotational experiment might prove to be less demanding and easier to visualize.  We conclude with some further perspectives on the analogy between viscoelastic and MHD flows.  The analogy is asymptotic in nature and therefore not perfect. Viscoelastic and MHD flows deviate from simple stress freezing in different ways: the viscoelastic stress relaxes, while the magnetic field diffuses.  The classes of exactly steady solutions of the two systems are therefore different, because after a sufficiently long time either relaxation or diffusion will have its effect.  For example, while there is only one solution for viscoelastic Couette flow, magnetized Couette flow can be set up with either vertical or azimuthal current-free magnetic fields.  In this sense, the analogy is more applicable to dynamical, time-dependent situations than to steady flows.  Renardy (1997) has analysed the large-${\\it De}$ limit of steady, two-dimensional flows of the upper-convected Maxwell fluid (obtained by setting the solvent viscosity $\\mu$ of the Oldroyd-B fluid to zero).  Noting that the stress $\\stress$ has one dominant eigenvalue in this limit, he writes $\\stress=\\rho\\bu\\bu$, where $\\rho$ and $\\bu$ are a fictitious density and a fictitious velocity field.  When inertial forces are negligible, $\\bnabla\\bcdot\\stress$ must balance the pressure gradient and $\\rho$ and $\\bu$ are then found to satisfy the steady Euler equations.  Through our analogy, a connection can be seen here with the work of Moffatt (1985), which makes use of the analogy between steady Euler flows and magnetostatic equilibria in which the Lorentz force balances the pressure gradient.  There is a close connection between the polymeric stress $\\stress_{\\rm p}$ at {\\it moderate\\/} De and the mean stress tensor $\\langle\\bB\\bB\\rangle/\\mu_0$ of a {\\it disordered\\/} magnetic field, as occurs in MHD turbulence.  Under these conditions, $\\stress_{\\rm p}$ and $\\langle\\bB\\bB\\rangle/\\mu_0$ may have three positive eigenvalues of comparable magnitude.  Indeed, Ogilvie (2001) has suggested the use of Maxwellian viscoelastic models, with Deborah numbers of order unity, for MHD turbulence in accretion discs.  Besides Couette flow, another problem that has received much attention is the stability of a planar jet or shear layer with respect to two-dimensional disturbances.  Azaiez \\& Homsy (1994) and Rallison \\& Hinch (1995) examined the equivalent problem for a viscoelastic fluid, noting the potentially stabilizing influence of a polymer additive. Taking a limit in which ${\\it Re}$ and ${\\it De}$ tend to infinity while maintaining a finite ratio, they derived an elastic equivalent of Rayleigh's stability equation.  An analogy exists between this problem and that of the stability of a similar flow of an ideal, electrically conducting fluid with a magnetic field parallel to the flow, a problem studied since the 1950s.  In a recent study, Hughes \\& Tobias (2001) derived a magnetic Rayleigh equation (their equation (3.5)) that is exactly equivalent to the elastic Rayleigh equation given by Rallison \\& Hinch (1995, p.~314).  In this paper we have drawn attention to a useful analogy between viscoelastic and MHD flows, and have discussed the relation between instabilities of differential rotation in the two systems.  We anticipate that much further use can be made of this analogy."
    },
    "0212/astro-ph0212169_arXiv.txt": {
        "abstract": "Structures on very large scales ($> 100 \\,\\rm Mpc$) have negligible peculiar motions, and are thus roughly fixed in comoving space. We looked for significant peaks at very large separation in the two-point correlation function --- corrected for redshift selection effects --- of a well convered subsample of 2378 quasars of the recently released 10k sample of the {\\sf 2dF} quasar survey. Dividing our sample in three redshift intervals, we find a peak at $\\simeq 244 \\, h^{-1} \\, \\rm Mpc$, which is perfectly comoving for a restricted set of cosmological parameters, namely $\\Omega_m = 0.25\\pm0.15$ and $\\Omega_\\Lambda=0.65\\pm0.35$ (both at 95\\% confidence). Assuming a flat Universe, we constrain the quintessence parameter $w_Q < -0.35$ (95\\% confidence). We discuss the compatibility of our analysis with possible peaks in the power spectrum. ",
        "introduction": "The quest for the parameters of the primordial Universe has significantly advanced in the last 10 years. There is a general agreement that $\\Omega_m \\simeq 0.3$, for example from the internal kinematics of clusters and groups \\cite{M93_aussois_dyn}, from the gas fraction in clusters \\cite{WNEF93} and groups \\cite{HM94}. Supernovae used as standard candles have led to $\\Omega_\\Lambda > 0$, with a degeneracy with $\\Omega_m$. Finally, the scales of the angular fluctuations of the cosmic microwave background (CMB) lead to a nearly perfectly flat Universe $\\Omega_m\\!+\\!\\Omega_\\Lambda = 1$. The combination of constraints from supernovae and the CMB recovers $\\Omega_m \\simeq 0.3$ (with $\\Omega_\\Lambda \\simeq 0.3$).  We present below a measurement of the cosmological parameters, with a fairly new technique based upon using very large scale structures ($> 100\\,\\rm Mpc$ in size) as comoving standard rulers. ",
        "conclusions": ""
    },
    "0212/astro-ph0212219_arXiv.txt": {
        "abstract": "Low-mass cluster galaxies are the most common galaxy type in the universe and are at a cornerstone of our understanding of galaxy formation, cluster luminosity functions, dark matter and the formation of large scale structure.  I describe in this summary recent observational results concerning the properties and likely origins of low-mass galaxies in clusters and the implications of these findings in broader galaxy formation issues. ",
        "introduction": "Although they are the faintest and lowest mass galaxies in the universe, low-mass cluster galaxies (LMCGs), especially dwarf ellipticals, hold clues for the ultimate understanding of galaxy formation, dark matter and structure formation. The reason for this is quite simple: low-mass galaxies, and particularly low-mass galaxies in clusters (Conselice et al. 2001) are the most common galaxies in the nearby universe (Ferguson \\& Binggeli 1994).  Any ultimate galaxy evolution/formation scenario must be able to predict and accurately describe the properties of these objects. In galaxy formation models, such as hierarchical assembly (e.g., Cole et al. 2000), massive dark halos form by the mergers of lower mass ones early in the universe.  By understanding these LMCGs, we are potentially studying the very first galaxies to form.   On the other hand, observations reveal that no low-mass galaxies formed all of their stars early in the universe at $z > 7$, with considerable evidence for star formation occurring in the last few Gyrs (e.g., Grebel 1997; Conselice et al. 2003).  While low-mass galaxies are traditionally studied in low density environments, such as in the Local Group, it is now clear that a large population of these low-mass galaxies exist in clusters, whose nature is only recently becoming clear (Conselice et al. 2001; 2003).  A comparison with the Local Group demonstrates that the ratio of low-mass to large mass galaxies in clusters is roughly five to ten times higher than in low density environments.  This over density of LMCGs, and the fact that some Local Group dwarf spheroidals (Klyena et al. 2002) have large dark matter halos, hints that potentially a large amount of mass in clusters is associated with low-mass galaxies.  New observational results also suggest that early-type LMCGs are not a homogeneous population, but consist of at least two distinct types, that are morphologically similar, but with different physical properties. ",
        "conclusions": ""
    },
    "0212/astro-ph0212505_arXiv.txt": {
        "abstract": "{We use recent population synthesis results to investigate the distribution of pulsars in the frequency space, having  a gravitational strain high enough to be detected by the future generations of laser beam interferometers. We find that until detectors become able to recover the entire population, the frequency distribution of the 'detectable' population will be very dependent on the detector  noise curve. Assuming a mean equatorial deformation $\\varepsilon =10^{-6}$, the optimal frequency is around 450 Hz for interferometers of the first generation (LIGO or VIRGO) and shifts toward 85 Hz for advanced detectors. An interesting result for future detection stategies is the significant narrowing of the distribution when improving the sensitivity: with an advanced detector, it is possible to have 90\\% of detection probability while exploring less than 20\\% of the parameter space (7.5\\% in the case of $\\varepsilon =10^{-5}$). In addition, we show that in most cases the spindown of 'detectable' pulsars represents a period shift of less than a tens of nanoseconds after one year of observation, making them  easier to follow in the frequency space. ",
        "introduction": "The coming years will be marked by the beginning of operations for the major gravitational wave interferometric antennas LIGO, VIRGO, GEO and TAMA. Starting with a sensitivity of the order of h $\\approx 10^{-22}$, they are expected to evolve in a few years into second generation experiments with a sensitivity improved by one or two orders of magnitude depending on the frequency. Ground based detectors will scan the sky searching for gravitational waves with frequencies between  tens of hertz up to tens of kilohertz. Potential sources fall roughly into three classes: bursts, stochastic background, continuous, involving different search techniques as match filtering for coalescing binary systems and cross-correlation between detectors for the stochastic background. The detection of pulsars (and other nearly monochromatic continuous sources) can be achieved by integrating the signal during times of about $10^{7}$ s and searching for statistically significant peaks at fixed frequencies in the power spectrum. This method becomes quite complicated as soon as one consider that: 1) pulsars are not really monochromatic emitters. The Doppler effect due to the Earth motion and the intrinsic spindown of the star become significant over long period of times spreading power across many frequency bins and must be corrected before performing the fourier transform. 2) most of pulsars are not observed in the electromagnetic domain due to strong selection effects, making necessary an all-sky and all-frequency search. The computational cost is therefore the main concern of searching for continuous sources: if the whole parameter space is scanned, the computing power required to process the data will be out of any reasonable value: of the order of $3 \\times 10^{20}$ Tera-flops for integration times of one year, frequency bandwiths of 2500Hz and decay times of 1000 year (Frasca et al. 2000). Alternative search algorithms involving hierarchical methods which follow up candidate detections from a first pass search have been investigated during the past few years (Brady et al. 2000, Schutz \\& Papa 1999, among others). These techniques reducing very significantly the computing time, by more than ten orders of magnitude for a small loss of sensitivity (by less than an order of magnitude) (Frasca et al. 2000) represent the best hope of detection for actual computational ressources and detector sensitivities. An attempt to optimize these algorithms could be the determination of the most likely range on the parameter space. Without restricting the search area definitively, the probability to have a first detection early could be optimized. Brady et al. (1998) have considered a survey in a limited sky area, searching for pulsars in the galactic core. In a previous work (Regimbau \\& de Freitas Pacheco 2000, hereafter paper I), population synthesis of galactic normal pulsars was performed in order to recover the statistical properties of the real population. This approach permits the modeling of the galactic pulsar contribution to the gravitational emission and to deduce its detectability by laser beam interferometers. The gravitational strain being proportional to the star equatorial deformation, the number of detection depends on the mean value adopted for the equatorial ellipticity. This parameter is quite uncertain since it depends on the cooling history and on different strains underwent by the crust. Past studies (Thorne 1980) have shown that equatorial distortions cannot be larger than $10^{-5}$, expected to be an upper limit related to the maximum stress that the solid crust could support, but it is worth mentioning that it could be several orders of magnitude smaller. Our simulations have showed that $10^{-6}$ is the critical value to have at least one detection with interferometers of the first generation ($10^{-7}$ for an advanced interferometer) In this paper, by using our previous population synthetis we investigate the detection probability of a pulsar in the frequency space. In order to illustrate our findings, we report here results for mean ellipticities equal to $10^{-5}$ and $10^{-6}$, as well as how the results are affected by the detector sensitivity. ",
        "conclusions": ""
    },
    "0212/astro-ph0212426_arXiv.txt": {
        "abstract": "We present new relations between recently defined line-strength indices in the near-IR (CaT$^*$, CaT, PaT, MgI, and sTiO) and central velocity dispersion ($\\sigma_0$) for a sample of 35 early-type galaxies, showing evidence for significant anti-correlations between Ca\\,{\\sc ii} triplet indices (CaT$^{*}$ and CaT) and $\\log \\sigma_0$. These relations are interpreted in the light of our recent evolutionary synthesis model predictions, suggesting the existence of important Ca underabundances with respect to Fe and/or an increase of the dwarf to giant stars ratio along the mass sequence of elliptical galaxies. ",
        "introduction": "During the last decade, the measurement and interpretation of blue optical line-strength indices in the spectra of early-type galaxies in the field have revealed the existence of an apparent spread of mean ages (Gonz\\'{a}lez 1993; Faber et al. 1995; J{\\o}rgensen 1999) and element abundances ratios (Worthey 1998; Trager et al. 2000a), suggesting a variety of interpretations of scaling relations like the colour--magnitude or Mg$_2$--$\\sigma$ relations (Bender, Burstein \\& Faber 1993; Kuntschner 2000; Trager et al. 2000b; Vazdekis et al. 2001). Since the picture from the blue is rather confused, if one wants to achieve a more complete understanding of the star formation history of these galaxies, it is necessary to look at other spectral regions in which the relative contribution of the distinct stellar types is very different. In this sense, the potential of the near-IR spectral range and, in particular, of the Ca\\,{\\sc ii} triplet is still almost unexploited.  Since Ca is an $\\alpha$-element like Mg, it should be enhanced compared to Fe in giant ellipticals (Es). However, as suggested by several authors, Ca seems to follow Fe (O'Connell 1976; Vazdekis et al. 1997; Worthey 1998; Moll\\'{a} \\& Garc\\'{\\i}a-Vargas 2000; Vazdekis et al. 2001; Proctor \\& Sansom 2002). Even so, given that variations of the Fe line-strengths among Es are not negligible (e.g. Gorgas, Efstathiou \\& Arag\\'{o}n-Salamanca 1990; Gonz\\'{a}lez 1993; Davies, Sadler \\& Peletier 1993; Kuntschner 2000), one should not expect the small variation of the Ca\\,{\\sc ii} triplet strength reported by previous work (Cohen 1979; Bica \\& Alloin 1987; Terlevich et al 1990; Houdashelt 1995). Furthermore, this result is difficult to interpret in the light of previous stellar population models (Garc\\'{\\i}a-Vargas et al. 1998; Schiavon et al. 2000) which predict a high sensitivity of the Ca\\,{\\sc ii} triplet to the metallicity of old, metal-rich stellar populations. Also, the absolute values of Ca\\,{\\sc ii} in Es differ from the model predictions (Peletier et al. 1999; Moll\\'{a} \\& Garc\\'{\\i}a-Vargas 2000).  With the aim of clarifying the above inconsistencies, during the last years we have developed a new stellar library in the near-IR spectral range (Cenarro et al. 2001a, hereafter CEN01) with a homogeneous set of revised atmospheric parameters for the library stars (Cenarro et al. 2001b), deriving empirical fitting functions that describe the behaviour of new line-strength indices for the Ca\\,{\\sc ii} triplet and the H Paschen series (CEN01; Cenarro et al. 2002) and other spectral features (Cenarro 2002, hereafter CEN02). Finally, in Vazdekis et al. (2002; hereafter VAZ02) we present a new evolutionary stellar population synthesis model which predicts both the integrated indices and the spectral energy distribution for single stellar populations (SSPs) of several ages, metallicities and initial mass functions (IMFs).  In this letter we present the first results for a spectroscopic sample of 35 early-type galaxies. After a brief description of observations and data reduction (Section~\\ref{stellib}), in Section~\\ref{indexsigma} we describe the measurements of the new indices for the central regions of the galaxies and their relationship with the velocity dispersion. In Section~\\ref{discussion} we discuss plausible interpretations of the data on the basis of new index--index diagrams derived from our model predictions. ",
        "conclusions": "\\label{discussion}  \\begin{figure*} \\centerline{\\hbox{ \\psfig{figure=fig02ref.epsi,width=17.5cm} }} \\caption{\\small CaT$^*$--sTiO and MgI--sTiO diagrams for the galaxy sample corrected to the system defined by the models at 370\\,km\\,s$^{-1}$ spectral resolution. Panels (a) and (b) show SSPs model predictions with fixed Salpeter IMF slope ($\\mu = 1.3$). Age varies from 5.01\\,Gyr (dotted line) to 17.78\\,Gyr (dash-dotted line) with solid lines for intermediate values ($\\Delta\\log$\\,[age(Gyr)] $\\sim 0.05$). Metallicity spans from $-0.68$ to $+0.20$ as in the labels (dashed lines). Panel (c) shows SSPs model predictions at fixed age of 17.78\\,Gyr with varying power-like IMF slope ($\\mu = 0.3 - 3.3$, see the labels) and metallicity (as in panel a). Different symbols indicate galaxies within distinct ranges of central velocity dispersion as it is shown in the key (b). Typical error bars for the whole sample are given.} \\label{3panels} \\end{figure*}  To analyse the previous relations we make use of our SSP model predictions (VAZ02; CEN02), which were transformed to the spectral resolution of the data (370\\,km\\,s$^{-1}$) using specific polynomials corresponding to their own age, metallicity and IMF (see VAZ02).  Given that the time evolution of the near-IR indices is virtually null for SSPs of all metallicities and ages $\\ga 3$\\,Gyr (VAZ02, CEN02), we can consider that the age has a negligible effect in this spectral range. Figure~\\ref{3panels} shows the distribution of the galaxy sample in the CaT$^*$--sTiO and MgI--sTiO planes, with symbols indicating different ranges of $\\sigma_0$. The insensitivity of the indices to age is apparent in Figs.~\\ref{3panels}a,b. From Fig.~\\ref{3panels}a, one immediately can notice that, while low-mass Es (filled circles) can be roughly fitted by with SSPs of metallicity below solar ($\\sim -0.4$\\,dex), no age-metallicity combination can account for the low CaT$^*$ values of massive Es (filled squares). The spread and location of the most massive Es (asterisks) will be further discussed. Moreover, typical random errors in CaT$^*$ are not able to explain the spread among Es, suggesting that other parameters must be taken into account to explain the observed trend. In this sense, we consider the existence of i) non-solar abundance ratios (hereafter NSAR), and/or ii) systematic variations of the IMF, as possible solutions to the above inconsistency.  NSAR could play an important role to interpret Fig.~\\ref{3panels}a. In fact, the high MgI values in massive Es (Fig.~\\ref{3panels}b) confirm the behaviour found for Mg in the optical. In this scenario, the position of massive Es in Fig.~\\ref{3panels}a implies that Ca should be underabundant in these galaxies, as was already suggested by e.g. Vazdekis et al. (1997) using the Ca4227 Lick/IDS index. As discussed in VAZ02, the use of $\\alpha$-enhanced scaled-solar isochrones (Salasnich et al. 2000) predicts even larger CaT$^*$ values (by $\\sim 0.5$\\AA) for metal rich, old SSPs. Therefore, given that overabundances of other $\\alpha$-elements exist, and the contribution of Ca to the total metallicity is lower than 0.5 percent, massive Es should have extremely low abundances of Ca to account for the observed values. Isochrones with accurate degrees of enhancement for the different $\\alpha$-elements as well as stellar spectral libraries with appropriate element ratios are indeed necessary to tackle with the above problem.  Could it be that massive Es are much more metal rich than the models, i.e., much more than two times solar metallicity? In this case, the shift of the red giant branch to lower temperatures would predict a significant decrease in the CaT$^*$ and CaT values. However, one would expect then sTiO values much larger than those observed (and also higher Fe values in the optical), but the measured indices are in reasonable agreement with the current model predictions.  The alternative to interpret the data on the basis of the current model predictions is to assume that the IMF may be varying among Es. Since age can be considered a secondary parameter, in Fig.~\\ref{3panels}c we plot the SSPs model predictions for 17.78\\,Gyr (a different age does not qualitatively alter the final conclusions) and different IMF slopes and metallicities. In this new scenario, the data can be interpreted by means of IMF slope and metallicity variations, suggesting that low-mass Es (filled circles) exhibit lower metallicities and lower IMF slopes (lower dwarf/giants ratio) than massive Es (filled squares). In the light of this interpretation, reading $\\mu$ and [Fe/H] values from the grid in Fig.~\\ref{3panels}c allows us to fit a one-parameter relation between $\\mu$, [Fe/H] and $\\log \\sigma_0$ for Es in the sense that, the larger $\\sigma_0$, the larger the metallicity and the IMF slope (Figure~\\ref{FundTube}). Note, however, that very massive Es (asterisks) deviate from the above relation exhibiting lower metallicities and slightly flatter IMFs. To guide the eye, solid lines represent simple polynomial fits to the data on the $\\mu$--[Fe/H] and [Fe/H]--$\\log \\sigma_0$ planes, whilst the relation in the $\\mu$--$\\log \\sigma_0$ plane is readily derived from the two previous ones. Using Monte-Carlo simulations (following a procedure similar to that of Kuntschner et al. 2001), we have checked that this relation is not driven by a combined effect of Poisson noise and non-orthogonal diagnostic diagrams, thus concluding that Es can indeed be explained using $\\mu$--[Fe/H]--$\\log \\sigma_0$ curves.  \\begin{figure} \\centerline{\\hbox{ \\psfig{figure=fig03ref.epsi,width=8.cm,angle=-90} }} \\caption{\\small The $\\mu$--[Fe/H]--$\\log \\sigma_0$ relation derived from Fig.~\\ref{3panels}c. Panels show the projection of the galaxies onto the planes $\\mu$--$\\log \\sigma_0$ (a), [Fe/H]--$\\log \\sigma_0$ (b), and $\\mu$--[Fe/H] (c). Lines indicate least-squares polynomial fits to all data: (b) ${\\rm [Fe/H]} = -1.09 + 0.46\\;\\log \\sigma_0$; (c) $\\mu = 2.41 + 2.78\\;{\\rm [Fe/H]} - 3.79\\;{\\rm [Fe/H]}^{2}$; (a) is straightly derived from (b) and (c). Symbols are as in Fig.~\\ref{3panels}.} \\label{FundTube} \\end{figure}  Although the universality of the IMF is a highly controversial question (see e.g. Eisenhauer 2001 for a review), there are theoretical arguments suggesting that the IMF of metal-rich star-forming regions must be biased towards low mass stars due to a more efficient cooling rate (e.g. Larson 1998). On this base, the universal IMF by Padoan et al. (1997) --which depends on the physical conditions of the star formation site-- predicts such a behaviour. The observational $\\mu$--[Fe/H] trend of Fig.~\\ref{FundTube}c qualitatively agrees with these theoretical arguments, but it still needs to be related with the galactic mass (or $\\log \\sigma_0$). A time-extended (non-instantaneous) star formation history or several star-forming episodes could drive such a relation. If mergers and accretion of other galaxies with pre-enriched gas have been frequent phenomena during the evolution of massive Es, subsequent star formation episodes should increase the metallicity and, as a consequence, the IMF would gradually tend to produce a larger relative number of low mass stars. In this case, the infall of non-primordial, previously enriched, gas is needed to ensure the increasing metallicity. Although an extended star formation could be enough to explain the increasing metallicity, the addition of dissipative hierarchical scenarios would emphasize such a behaviour.  This interpretation implicitly characterizes the IMF as a time-dependent function (it gets steeper with time) in agreement with several previous work (e.g. Larson 1998), suggesting that the IMF in the early Universe could be biased towards very massive stars. A prompt initial enrichment --straight consequence of a time-dependent IMF-- could account for the observed $\\alpha$-enhancements (Vazdekis et al. 1996). This could also explain several keys like the G-dwarf problem, the scarcity of Population III stars, and the high abundance of heavy elements in the intergalactic medium of clusters of galaxies (Larson 1998). In addition, the existence of a $\\mu$--[Fe/H]--$\\log \\sigma_0$ relation would also have important consequences for the interpretation of the blue spectra of Es. In particular, the $<$Fe$>$--$\\log \\sigma_0$, H$\\beta$--$\\log \\sigma_0$, and the mass--luminosity (in $B$ band) relations could only be reconciled by introducing an age sequence in the sense that the larger $\\log \\sigma_0$ the lower the mean age in the central parts (CEN02). Another problem could be visual-infrared colors. For example, for a model with $\\mu = 2.8$, [Fe/H] $= +0.2$ and 12.6\\,Gyr one expects $V-K = 3.52$ (Blakeslee et al. 2001).  This is just at the edge of the range of observational values for Es (e.g., $V-K \\la 3.50 \\pm 0.05$ for Es in Frogel et al. 1978). However, and again, a younger mean age for the central regions would help to match the observed values ($V-K = 3.27$ for a model with $\\mu = 2.8$, [Fe/H] $= +0.2$ and 5.01\\,Gyr).  The interpretation with a varying IMF actually revisits the classic debate about the existence of a dwarf-enriched population in the nuclei of Es, based on the strengths of the near-IR Na\\,{\\sc i} doublet, the Ca\\,{\\sc ii} triplet and the FeH Wing-Ford band (Cohen 1978; Faber \\& French 1980; Carter et al. 1986; Alloin \\& Bica 1989; Couture \\& Hardy 1993). In particular, the weakness of the FeH band was used as an argument against this possibility. Since these results are mainly based on rather limited empirical synthesis models, they are not strong enough to directly exclude the possibility of a dwarf-heavy IMF. A proper calibration of its sensitivity to the stellar parameters (in particular to metallicity; see Carter et al. 1986) and its inclusion in modern stellar population models must be performed before extracting any conclusion. In any case, since several model uncertainties affect the absolute scale of the predicted index strengths (by $\\sim$ 0.5\\,\\AA\\ in the case of CaT$^*$; see VAZ02), the derived $\\mu$--[Fe/H]--$\\log \\sigma_0$ relation must be considered on a relative basis.   Finally, we speculate on the very massive Es labelled with asterisks in Figs.~\\ref{3panels} and \\ref{FundTube}. Whereas all (except one) are boxy Es, with slow rotation and resolved cores, the rest of the sample are mostly disky Es, fast rotators with power-law cores. According to Faber et al. (1997), disky Es are consistent with their formation in gas-rich mergers, whereas boxy Es could be the by-products of gas-free stellar mergers. In the context of our interpretation, this suggests that boxy Es should exhibit lower metallicities (and, therefore, lower sTiO values) at the time that keep a more primordial, flatter slope IMF than disky Es (thus leading to larger values of CaT$^*$ and CaT), in agreement with what we observe in Figs.~\\ref{I-sigma} and ~\\ref{FundTube}. Also, the fact that boxy Es seem to be older than disky Es (de Jong \\& Davies 1997; Ryden, Forbes \\& Terlevich 2001) favours the last interpretation.  We want to note that, at the time of submission of this letter, Saglia et al. (2002) presented similar anti-correlations between Ca\\,{\\sc ii} indices and the velocity dispersion, using another large sample of high-quality spectra. Another paper by Falc\\'on-Barroso et al. (2003) has been submitted with such anti-correlations for bulges of spiral galaxies.  To conclude, more work in the areas of SN yields, stellar interior models and stellar libraries is needed to clarify whether it is possible to explain the current discrepancy between Ca\\,{\\sc ii} measurements and models using non-solar abundance ratios. Until this is accomplished, and in the light of the present models, the only way out is to advocate for variations in the IMF. By no means we think that we are presenting strong evidence against its universality. Furthermore, the scenario of a time-dependent IMF poses difficulties to explain other observables. More modern, careful work using indicators in the near-IR and the optical is needed to resolve the issue of a non-standard IMF."
    },
    "0212/astro-ph0212083_arXiv.txt": {
        "abstract": "\\noindent The braneworld model of Dvali--Gabadadze--Porrati (DGP) is a theory where gravity is modified at large distances by the arrested leakage of gravitons off our four-dimensional universe.  Cosmology in this model has been shown to support both ``conventional\" and exotic explanations of the dark energy responsible for today's cosmic acceleration.  We present new results for the gravitational field of a clustered matter source on the background of an accelerating universe in DGP braneworld gravity, and articulate how these results differ from those of general relativity.  In particular, we show that orbits nearby a mass source suffer a universal anomalous precession as large as $\\pm 5~\\mu{\\rm as/year}$, dependent only on the graviton's effective linewidth and the global geometry of the full, five-dimensional universe.  Thus, this theory offers a local gravity correction sensitive to factors that dictate cosmological history. ",
        "introduction": "The gravity theory of Dvali--Gabadadze--Porrati (DGP) is a braneworld theory with a metastable four-dimensional graviton \\cite{Dvali:2000hr}.  The graviton is pinned to a four-dimensional braneworld by intrinsic curvature terms induced by quantum matter fluctuations; but as it propagates over large distances, the graviton eventually evaporates off the brane into an infinite volume, five-dimensional Minkowski bulk.  There exists a single free parameter in DGP braneworld gravity, the crossover scale, $r_0$.  This scale dictates that distance below which gravity is controlled by brane effects, but larger than which gravity assumes a five-dimensional behavior.  As a result, the DGP braneworld theory is a model in a class of theories in which gravity deviates from conventional Einstein gravity not at short distances (as in more familiar braneworld theories), but rather at long distances.  Such a model has both intriguing phenomenological \\cite{Dvali:2001gm,Dvali:2001gx} as well as cosmological consequences \\cite{Deffayet,Deffayet:2001pu,Deffayet:2002sp,Alcaniz:2002qh,Jain:2002di,Alcaniz:2002qm,Lue:2002fe}. A braneworld model of the sort where gravity is modified at extremely large scales is motivated by the desire to ascertain how our understanding of cosmology may be refined by the presence of extra dimensions.  Indeed, there exist novel cosmologies in this theory that provide an alternative explanation of the cosmic dark energy \\cite{Deffayet}.  Recovery of Einstein gravity at short distance scales in DGP \\cite{Deffayet:2001uk}, especially for static sources \\cite{Lue:2001gc,Gruzinov:2001hp,Porrati:2002cp,Middleton:2002qa}, is a subtle effect.  Even though gravity is four-dimensional at distances less than $r_0$, it is not always Einstein gravity. For a point source whose Schwarzschild radius is $r_g$, general relativity is only recovered for distances shorter than \\be r_* = \\left(r_0^2r_g\\right)^{1/3}\\ . \\ee For distances larger than $r_*$, gravity is four-dimensional linearized Brans--Dicke, with parameter $\\omega = 0$.  Thus, a marked departure from conventional physics continues down to distances much smaller than $r_0$, the distance at which the extra dimension is naively hidden.  That departure is well-characterized and provides a possible signature for the existence of extra dimensions.  There is, however, a catch.  The cosmological solutions that drive interest in DGP gravity indicate that $r_0$ should be close to today's Hubble radius.  Localized matter sources are embedded in a cosmological spacetime.  Far enough away from a given source, the motion of observers and test particles is dominated by the cosmology, rather than the gravity of the matter source.  However, if $r_0$ is today's Hubble radius, then the distance away from a source at which cosmology dominates the metric is also $r_*$.  So, the DGP departure from Einstein gravity that was calculated in a static Minkowski background is only significant on length scales where the background cannot actually be approximated as Minkowski.  The calculations needs to be refined if we are to have the correct new physics signature.  The subject of this paper is to look at corrections to Einstein gravity for the DGP braneworld theory in a cosmological background.\\footnote{ The work in this paper, as well as that in Refs.~\\cite{Lue:2001gc,Gruzinov:2001hp,Porrati:2002cp}, addresses issues similar to those studied in Refs.~\\cite{Kofinas:2001qd,Kofinas:2002gq}.  Our results differ from those in the latter work because we are only concerned with solutions that asymptotically approach well-behaved, familiar solutions (e.g., Minkowski space or deSitter expansion) far away from the matter source, including far away from the source in the direction into the bulk.} We begin by laying out the necessary details of the model as well as the particular background of interest.  We then solve for the metric of a spherically symmetric, static matter source in that background. We indeed find that the corrections to Einstein are sensitive to the background cosmology, even in the region inside $r_*$ where the cosmological flow is ostensibly irrelevant.  We discuss possible tests for the detection of new physics at astronomical scales and suggest that it is possible for such tests of local gravity to reveal information about global features of the full five-dimensional cosmology, and ultimately shed light on the nature of dark energy and today's cosmic expansion. ",
        "conclusions": "The braneworld theory of Dvali--Gabadadze--Porrati (DGP) is an intriguing extension of Einstein gravity that exploits the possible existence of infinite-volume, extra dimensions.  It is a theory where the four-dimensional graviton is effectively metastable, and provides a novel alternative to conventional explanations of the dark energy that is responsible for today's cosmic acceleration.  In this paper, we detailed the solution of static, spherical matter sources in the background of deSitter cosmology for DGP gravity. The gravitational field of a matter source exhibits important dependences on cosmology.  Residual dependences on the full five-dimensional cosmological phase also exist in the regime deep in the gravity well of the matter source where the effects of cosmology are ostensibly irrelevant.  These residual dependences allow one to use local (e.g., solar system) measurements of the gravitational field to ascertain details of the global cosmology.  In DGP gravity, we find that massless test particles can probe the true mass of a matter source, whereas tests of the source's Newtonian force leads to discrepancies with general relativity. These discrepancies translate into a universal anomalous precession, as large as $\\pm 5~\\mu{\\rm as/year}$, suffered by all orbiting bodies.  The numerical value of this anomalous precession is dependent only on the graviton's effective linewidth and the global geometry of the five-dimensional cosmology.  Current constraints on Mars' orbit are on the threshold of being sensitive to this anomaly \\cite{Nordtvedt:ts}.  Future improvements in lunar ranging \\cite{lunar} as well as data from satellite missions at the end of the decade \\cite{Milani:2002hw} should be sensitive to possible corrections due to DGP and gravitational leakage into extra dimensions."
    },
    "0212/astro-ph0212206_arXiv.txt": {
        "abstract": "{ We are exploring the  connection between damped Ly-alpha absorbers (DLAs) s and Lyman break galaxies (LBGs) using deep -- m$_{lim,I}$(5$\\sigma$)=26 m$_{AB}$-- broad band imaging (UBVI) of four wide fields (0.25deg$^2$ each)  obtained  at the Kitt Peak 4-m telescope with MOSAIC. Each field contains a DLA at $z\\sim3$.  We want to address the nature of DLAs at high-redshifts: (1)  Are they embedded in much larger systems of galaxies? (2) How does the spatial distribution of LBGs in 3D (space and redshift) correlate with the absorber? Contrary to most previous DLA studies, we are not looking for the absorber, and we do not rely on control fields because each of our fields is  $40\\times 40 h^{-1}$ Mpc (co-moving). We present preliminary results in two of our fields. In one case, we see an indication of an overdensity of galaxies on a scale of 5 Mpc. We discuss the possible implications and sources of contamination of our results.  } ",
        "introduction": "\\label{sec:intro} By definition, damped Ly-$\\alpha$ absorbers (DLAs) found in the spectra of quasars have neutral hydrogen (HI) column densities greater than $2\\times 10^{20}$ cm$^{-2}$. Their nature is unknown and has been an ongoing debate for more than a decade. Two classes of hypothesis can explain DLA properties, (1) DLAs are large disks or proto-disks, or (2) DLAs arise from the superposition of small gas clouds in a large halo. The former was first proposed by Wolfe et al. (1986), the latter is more recent. Maller et al. (2000) showed that DLAs can arise from the combined effects of massive central galaxies and a number of smaller satellite galaxies in a virialized halo. In this scenario, a DLA would lie in an over dense region.  Lyman Break Galaxies (LBGs) are actively star forming objects, hence are bright $L_{UV}=10\\times L^*(today)$ (e.g. Steidel 2000)  and they trace the underlying matter distribution as indicated by their strong clustering.  Based on this, the motivation of this work is to try to answer the following question: {\\it Are DLAs in over- or under-dense regions?} In order to test this, we selected four QSO fields that have (i) a low galactic extinction, and (ii) $z_{DLA}\\sim 3 < z_{QSO}$. In order to constrain the environment and the nature of the damped systems at $z\\sim3$, we imaged four fields with the MOSAIC camera at the Kitt Peak 4-m telescope in UBVI. The area covered (0.25deg$^2$/field to $I_{AB}(5\\sigma)<26$) allow us to identify several hundred LBG candidates both close to the QSO/DLA line of sight (50 kpc) and up to 20 $h^{-1}$ Mpc (co-moving) away. Thus, unlike previous studies, we do not rely on a randomly selected control field.  In this proceeding, we present preliminary results in the fields APM 08279+5255 and PC1233+4752 that contain a DLA at $z_{DLA}=2.97$ and $z_{DLA}=3.5$ respectively.  At $z\\sim3$,  $1\\arcmin$ corresponds to $\\sim 1.3h^{-1}_{100}$~Mpc (co-moving) for  $\\Omega_M=0.3$, $\\Omega_{\\Lambda}=0.7$. ",
        "conclusions": "\\label{sec:spatial} We selected high-redshift candidates to be U-band drop-outs with $21<I<24.5$. We then estimated the redshift according to section~\\ref{sec:photoz}. Our results, the spatial and radial distributions of LBGs candidates around the DLA, are shown on Figure~1 for APM08279+5255 separated in  two redshift bins around the DLA, $|z_{DLA}-z|<0.2$, and $0.2<|z_{DLA}-z|<0.5$. The surface density indicates an overdensity. The upper limit is 0.5 to account for incompleteness in our redshift distribution for our selection function drops sharply at $z\\sim2.8$. In each case, the radial surface density is plotted. Note that we find the same signature even when we split the first bin in half. Around PC1233+4752, we do not find such a signature.  Our results indicate that at least some DLAs may lie in overdense regions. We note that Adelberger et al. (2002) found  an under-density (marginally), and Gawiser et al. (2001) were inconclusive. Our results are potentially limited by our lack of accurate redshifts --- $\\Delta z= 0.1$ corresponds to 75 $h^{-1}$ Mpc at $z=3$---  and contamination of low redshift objects and red stars (estimated to be  10\\% or less). This would tend to reduce the overdensity signature. If our results were due to a clustering of stars, the signature would disappear when we increase our lower magnitude limit. We find the over-dense signature remains.  This Fall, we will obtain multi-object spectroscopy (with GMOS) of some of our candidates in order to test our photometric redshifts and to confirm our results. We are in the process of applying the same techniques to our two other fields (Bouch\\'e et al 2003).    \\begin{figure} \\begin{center} \\includegraphics[width=\\columnwidth]{nbouche_fig1.eps} \\end{center} \\caption{Left: Spatial distribution of LBG candidates. Right: Radial surface density as a function of radial distance from the QSO line-of-sight (shown with the cross). The lower panels are for the candidates with $|z_{DLA}-z|<0.2$, and the upper panels are for $0.2<|z_{DLA}-z|<0.5$. The dashed line is the average of density of the last three radial bins, representing the underlying density. A $1\\sigma$ signature of an overdensity is seen. \\label{fig1}} \\end{figure}"
    },
    "0212/astro-ph0212030_arXiv.txt": {
        "abstract": "{  We discuss multiple Very Long Baseline Interferometry (VLBI) continuum and spectral line imaging observations and Westerbork Synthesis Radio Telescope spectroscopy of the compact variable nuclear radio jet source in the elliptical galaxy \\object{NGC\\,1052}. Absorption and emission signatures reveal ionised, atomic, and molecular components of the surrounding medium.  Ten epochs of Very Long Baseline Array (VLBA) data at 15\\,GHz, spanning almost six years, show bi-symmetric jets, in which multiple sub-parsec scale features display outward motions of typically $v_{\\rm app}\\sim0.26c$ (${\\rm H}_0=65$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$) on each side. The jets are most likely oriented near the plane of the sky.  Multi-frequency VLBA observations at seven frequencies between 43 and 1.4\\,GHz show free-free absorption in the inner parsec around the nucleus, probably together with synchrotron self-absorption. The free-free absorption is apparently due to a structure which is geometrically thick and oriented roughly orthogonal to the jets, but which is patchy. The western jet is covered more deeply and extensively, and hence is receding.  H{\\sc i} spectral line VLBI observations reveal atomic gas in front of the approaching as well as the receding jet. There appear to be three velocity systems.  Broad, shallow absorption asymmetrically straddles the systemic velocity spanning $-35$ to 85\\,km\\,s$^{-1}$. This gas could be local to the AGN environment, or distributed on galactic scales. Superimposed in the range 25 to 95\\,km\\,s$^{-1}$ are several sharper (3--15\\,km\\,s$^{-1}$) features, each detectable over a few tenths of a pc at various places along the inner 2\\,pc of the approaching jet. The third, deepest system is at ``high velocities\", which is receding by 125 to 200\\,km\\,s$^{-1}$ with respect to the systemic velocity of \\object{NGC\\,1052}.  It may have a continuous velocity gradient across the nucleus of some 10\\,km\\,s$^{-1}$\\,pc$^{-1}$. This atomic gas seems restricted to a shell 1--2\\,pc away from the core, within which it might be largely ionised.  Westerbork Synthesis Radio Telescope spectroscopy has revealed the 18\\,cm OH main lines (1667 and 1665\\,MHz) in absorption along the full velocity span of $-35$ to 200\\,km\\,s$^{-1}$, with their line ratio varying roughly from 1:1 to 2:1. They are deepest in the high velocity system, where the OH profiles are similar to H{\\sc i}, suggesting co-location of that atomic and molecular gas, and leaving unclear the connection to the H$_2$O masing gas seen elsewhere.  In the high velocity system we have also detected the 18\\,cm OH satellite lines: 1612\\,MHz in absorption, and 1720\\,MHz in emission. The conjugate behaviour of the satellite line profiles, and the variable main line ratio resemble the situation in \\object{Cen~A} and \\object{NGC\\,253}. ",
        "introduction": "}  The elliptical galaxy \\object{NGC\\,1052} is a low luminosity AGN which has a well-studied LINER optical spectrum (e.g., \\gabpp) and an unusually prominent central radio source, with a flux density of 1--2\\,Jy and a fairly flat spectrum between 1--30\\,GHz, which is variable on timescales of months to years (e.g., \\heexx; \\heeie). Since \\object{NGC\\,1052} is comparatively nearby\\footnote{We adopt $H_0$=65\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, with no corrections for local deviations from the Hubble flow, so that the optical stellar absorption line heliocentric redshift, $cz=1474$\\,km\\,s$^{-1}$ or $z=0.0049$ (\\saruu), corresponds to a distance of 22\\,Mpc and implies that 1\\,mas corresponds to 0.1\\,pc}, very detailed scrutiny of the active nucleus and its inner environment are possible. This is important because the extension to lower luminosity of the active nucleus/relativistic jet phenomenon has not been well studied.  A 1.4\\,GHz Very Large Array (VLA) image of \\object{NGC\\,1052} shows a core-dominated radio structure, with only about 15\\%\\ of the flux density in extended emission: there are two lobes spanning $\\sim3$\\,kpc, possibly with hot spots (Wrobel \\cite{wro84}). The core luminosity at 5\\,GHz is $1.3\\times10^{23}$\\,W\\,Hz$^{-1}$ (\\keloi, hereafter K98). Thus, while \\object{NGC\\,1052} is clearly in the luminosity class of classical FR\\,I radio sources, its radio spectrum, morphology, and size (the entire radio source has dimensions smaller than the host galaxy), suggest that it may be akin to Gigahertz-Peaked-Spectrum or Compact-Steep-Spectrum sources, which are generally thought to be young jets, propagating through, and perhaps interacting with, a rich inner galactic medium (see the review by O'Dea \\cite{ode98}).  Early Very Long Baseline Interferometry (VLBI) observations showed that most of the radio emission is from a sub-pc scale core (\\cohuq), extended over $\\sim$15\\,milliarcseconds (mas) in position angle 63\\deg\\ (\\jonir). Recent work by Kadler et al.\\ (\\cite{kad02a}) shows that coincident X-ray, optical, and radio knots, at distances of some 500\\,pc on both sides of the core, are aligned with the parsec-scale radio jets. The emission extended further out to a few kpc seems to be more along a position angle of 95\\deg, as already seen in the VLA image by Wrobel (\\cite{wro84}). \\daviyp\\ found the parsec-scale jets to be approximately perpendicular to the gas rotation axis, which they measured to be at $-41$\\deg\\ using the optical emission lines. The same orientation is also seen in the neutral, H{\\sc i} 21\\,cm line emitting gas component (\\vaniy, hereafter vG86). But, as is not uncommon in elliptical galaxies, the gas rotation axis is oriented well away from that of the stars, which is at position angle 25\\deg, close to the projected photographic minor axis of the galaxy (e.g, \\daviy). \\object{NGC\\,1052} is probably a triaxial system, and it is plausible that the gas, which has sub-solar abundances (Sil'chenko \\cite{sil95}) and a gas-to-dust ratio comparable to that in our local Galaxy, was acquired as a result of one or more mergers or tidal interactions, for example with the nearby spiral \\object{NGC\\,1042} (see vG86; \\daviy; \\plaoi).  There is much observational evidence for the presence of a substantial amount of gas and dust along the line of sight to the nucleus. First, Chandra observations reveal an X-ray jet embedded into a 0.5\\,keV thermal plasma (Kadler et al.\\ \\cite{kad02a}).  Optical images show a faint diffuse kpc-scale dust lane extending along the projected minor axis (e.g., \\carie; \\daviy). Also, \\gabppp\\ have imaged an H$\\alpha$ emission line filament on scales of $\\sim$100\\,pc. Towards the nucleus there is reddening of the continuum light (e.g., \\carie), as well as the nebular emission lines (\\gabpp). Polarised optical broad lines (\\baroo) give evidence for the presence along our direct line of sight to the AGN of the kind of toroidal obscuring region often invoked in unification models for various kinds of active galaxies (e.g., \\antoe). The soft X-ray spectrum also shows strong absorption, indicating a substantial, and possibly inhomogeneous column depth of ionised gas towards the active nucleus (\\guaoo; \\guapp; \\weaoo; Kadler et al. \\cite{kad02a}).  H{\\sc i} 21\\,cm line absorption has been found towards the bright nuclear radio source at velocities ranging from $\\sim$1480--1680\\,km\\,s$^{-1}$, i.e.\\ extending redward from systemic over 200\\,km\\,s$^{-1}$ (\\shoie; vG86). Around 1640\\,km\\,s$^{-1}$ H$_2$O maser emission has been detected (\\braor) and imaged with the VLBA (\\claoi). Recently, 1667 and 1665\\,MHz OH absorption near 1640\\,km\\,s$^{-1}$ has also been detected (\\omapw).  The presence of CO lines is unclear: \\wanowp\\ claimed a detection of CO(1--0) in emission near 1420\\,km\\,s$^{-1}$, but the peak strength of 150\\,mK is much above the 20 mK 3$\\sigma$ upper limit claimed by \\wikotp, and is also not confirmed by \\knaoyp. Meanwhile, \\knaoyp\\ did find tentative absorption features in CO(1--0) in the velocity range 1550--1700\\,km\\,s$^{-1}$. As they point out, lines at that velocity might also be visible in the CO(1--0) and CO(2--1) spectra of \\weaotp. Since there is now confirmed molecular gas at this velocity in OH and also in H$_2$O masers, it is important to verify the CO results.  This paper deals with various aspects of the radio jets and the measurements of ionised atomic and molecular gas in their surrounding medium.  Section~\\ref{sec:obs} collects the details of all of the observations and the data reduction. In Sect.~\\ref{sec:kinematics} we discuss the morphology and kinematics of sub-pc scale features in the jets, based on 15\\,GHz VLBA observations made at a total of ten different epochs, spanning a period of more than five years. In Sect.~\\ref{sec:iongas} we present the broad-band continuum spectra of various jet components, which indicate extensive, asymmetric free-free absorption; this is based on VLBA observations at seven frequencies, spanning the range from 1.4\\,GHz up to 43\\,GHz. Kellermann et al.\\ (\\cite{kel99}, hereafter K99) reported on motions and free-free absorption in \\object{NGC\\,1052}, now fully described in this paper. Kameno et al.\\ (\\cite{kam01}, hereafter Ka01) later analysed another dataset also showing free-free absorption.  Kadler et al.\\ (\\cite{kad02a}; \\cite{kad02b}) analysed further multi-frequency data finding free-free absorption also, and linearly polarised emission at 5\\,GHz.  Section~\\ref{sec:atogas} shows the spatial and velocity distribution of the H{\\sc i} absorption on sub-pc scales, based on H{\\sc i} 21\\,cm spectral line VLBI observations. Section~\\ref{sec:molgas} deals with Westerbork Synthesis Radio Telescope (WSRT) observations of all four 18\\,cm OH lines, using the new wide band multi-channel digital backend (DZB) which is very well suited for observing spectral lines over wide bandwidths with multiple spectral channels.  We find the OH 1665 and 1667\\,MHz lines over a much wider velocity range than originally discovered by \\omapwp, and we also find for the first time the presence of the OH 1612 and 1720\\,MHz lines, in absorption and emission, respectively. ",
        "conclusions": "}  The general picture which emerges from the data presented in the previous sections supports the scenario often invoked in AGN (unification) models. The central engine produces twin radio jets, which lie along the polar axis of a toroidal or disk-like region into which matter is accumulating. However, some important questions remain.  The jet velocity in NGC\\,1052 is only mildly relativistic, rather than approaching the velocity of light, as in many other AGN.  Our observations may refer only to a slowly moving external surface of the jets, so that the radiation from a faster moving central spine is undetectable because it is beamed away from the line of sight. However, NGC\\,1052 is one of the lowest luminosity galaxies in which a jet velocity has been measured to date. It is quite possible that low luminosity objects intrinsically have slower jets on average.  The combination of free-free absorption, from a geometrically thick, albeit patchy, ionised structure covering the inner parsec, and H{\\sc i} atomic gas predominantly in an annulus 1--2\\,pc around the centre, is quite natural in principle. The innermost region and/or the surface of the accretion disk or torus receive the most intense radiation from the AGN, and thus have the largest degree of ionisation. Indeed, recent work by \\weaoop, \\guaoop, \\guappp, and Kadler et al.\\ (\\cite{kad02a}) has shown the presence of large column depths ($N_{\\rm H}$=10$^{22}$--10$^{24}$\\,cm$^{-2}$) of soft X-ray absorbing gas along our line of sight to the nuclear continuum source. Furthermore, some models of the X-ray spectral details suggest that the distribution of the absorbing gas is patchy, or that two components with a substantially different density are involved, matching the evidence we see for patchiness also in the radio spectra and flux density evolution of different jet components at various distances from the core. The detection by \\beaoop\\ of a broad H$\\alpha$ emission line, visible only in polarised (scattered) light, is also evidence for obscuration by a canonical equatorial torus or disk. But several inconsistent reports on the (non)detection of CO emission and absorption (\\weaot; \\wikot; \\knaoy) have left the amount of molecular gas near the nucleus uncertain.  However, the nature of the high velocity H{\\sc i} absorption system is not clear. The apparent 10\\,km\\,s$^{-1}$\\,pc$^{-1}$ velocity gradient could correspond to a rotation component of the motion. Naively, for matter at a distance of 1\\,pc from the core, this rotation would imply a modest enclosed mass of about 10$^7$\\,M$_{\\odot}$, but we suspect that the actual geometry is rather more complex. In particular, this system is redshifted by 130--200\\,km\\,s$^{-1}$ compared to the stellar recession velocity of NGC\\,1052. Since the spatial distribution strongly suggests that the gas is quite local to the AGN environment, it is most likely fairly cohesively falling in to the nucleus. One might speculate that this is the result of recent accretion of gas due to (a series of) mergers.  NGC\\,1052 has complex, probably triaxial kinematics on arcsecond scales, as revealed by optical spectroscopy (\\plaoi) and radio observations of H{\\sc i} (vG86), and it is not obvious how these large-scale motions connect to those of parsec-scale features. Under an inflow scenario, we are puzzled by the absence of H{\\sc i} absorption towards the innermost approaching jet, given that the three dimensional distance to the ionising central source may then not differ very much from that of gas seen 1--2\\,pc away from the nucleus in projection on either side. We point out, incidentally, that the radio jets do not appear to be ``significantly'' oriented with respect to the isophotal or rotation axes of either the stellar body or the neutral or ionised gas in NGC\\,1052, despite early suggestions to the contrary (\\jeair; \\daviy; vG86). While the initial direction of ejection is most likely determined by the rotation of the black hole, it is unclear what determines the subsequent bending.  While the WSRT observations of the 18\\,cm OH lines in NGC\\,1052 do not resolve the arcsecond-scale source, we believe that it is likely that in the high velocity system the OH lines are due to molecules situated towards the eastern, approaching jet, at distances of 1--2\\,pc from the core, just like the H{\\sc i} line. We know from our 1997 18\\,cm continuum VLBA data that the receding jet contained a flux density of only $\\sim$20\\,mJy, so that the the local OH opacity in the 1667\\,MHz line would have to be as high as 25\\%\\ if the absorption were to occur only on that side. Just as importantly, the high velocity part of the OH profile matches that of the integrated H{\\sc i} absorption rather well (Fig.~\\ref{ohall}). The peak H{\\sc i} absorption depth is 5\\% of the total flux density, as compared to 0.4\\%\\ in OH. Remembering that this is based on values integrated over the source, and with a 3.5 year time difference, we find that the opacity ratio of 12:1, which converts to a column depth ratio of $10^{6}$ (naively taking $T_{\\rm sp}=100$\\,K for H{\\sc i} and $T_{\\rm ex}=10$\\,K for OH), matches very well the models of \\meaoyp\\ for conditions near an AGN.  While we cannot rule out that the ``low velocity'' H{\\sc i} and OH absorption is the result of a long path length through the galaxy, the OH excitation points toward an association of this gas with the nuclear region.  The sharper, localised H{\\sc i} absorption lines also have no clear counterpart in OH, although the spectral resolution in OH was a factor 2--4 lower. Perhaps these are clouds with a different ionisation parameter and a lower fraction of molecular gas.  It is also possible that there are at least a few differently distributed molecular gas components present in the circumnuclear medium. \\claoip\\ have found H$_2$O maser emission, indicating the presence of molecular gas at 0.1--0.2\\,pc to the west of the core, along the receding jet (note that the zero point in the images of \\claoip\\ is not on the core, since they did not have the benefit of multi-epoch and multi-frequency datasets). Possibly, these masers are excited by interaction with the jets. We have no explanation, however, for the fact that while the H$_2$O masers have the same speeds as the ``high velocity system'', they appear to have a different location than the majority of the H{\\sc i} and OH line producing gas.  Remarkably, therefore, in the complex accretion region, which the inner parsec around the AGN surely is, it appears possible for some molecular and atomic clouds to preserve a relatively quiescent existence, with apparently highly ordered motions of the high velocity system, and kinematically fairly undisturbed separate sharp H{\\sc i} absorbing clouds. There are other active galaxies in which somewhat similar obscuring regions have been found, but they each seem to have their own peculiarities. For example, \\kaspwp\\ find at least three components with different ionisation parameters in X-ray absorption towards NGC\\,5548, with variable broad lines showing a location between 10$^{14}$ and 10$^{16}$\\,cm. In NGC\\,1275 VLBI data show that the approaching jet is apparently not covered by free-free absorption, but the inner portion of the counter-jet is hidden behind a thick or warped disk (e.g. \\walpp; \\levot). In NGC\\,4261 \\jonpqp\\ find evidence for a geometrically thin obscuring disk. There is also H{\\sc i} absorption in NGC\\,4261, but in this case, \\vanppp\\ have found that it occurs only towards the counter-jet, and that the column depth is only $N_{\\rm H}=2.5\\times10^{19}$\\,cm$^{-2}$. As a final example, in NGC\\,5793, \\pihppp\\ find that there are H{\\sc i} absorption components at several velocities, which each cover a different VLBI continuum component, and they span a region of 30\\,pc, rather larger than in NGC\\,1052. It is also intriguing to extend the comparison with NGC\\,5793 to molecular gas. The 1667 and 1665\\,MHz OH main lines have been imaged in absorption by \\hagppp.  Their extent is somewhat uncertain, but is at least 5\\,pc. Interestingly, there is a 9\\,km\\,s$^{-1}$\\,pc$^{-1}$ velocity gradient across the source, but in this case it is centred on $v_{\\rm sys}$. There is also H$_2$O maser emission in NGC\\,5793, but at another position and a blue-shifted velocity (\\diapq).  }  We have presented and discussed multiple VLBA continuum and spectral line imaging and WSRT spectroscopic observations of the compact variable nuclear radio source in the nearby active galaxy NGC\\,1052. Our main conclusions can be summarised as follows:  1. The radio source consists of slightly curved bi-symmetric jets with multiple sub-parsec scale features.  2. There are outward motions, reasonably linear on the sky and constant over time, of typically $v_{\\rm app}\\sim0.26c$ (${\\rm H}_0=65$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$) on each side.  3. Symmetry shows that the jets are oriented near the plane of the sky.  4. Spectral shapes of jet components ranging from steep, through convex, to highly inverted suggest that synchrotron self-absorption and free-free absorption are both at play along the inner parsecs of the jets.  6. Free-free absorption leads to an asymmetric central gap opening up towards lower frequencies. Thus, the western jet is receding, and is partially obscured along at least 1\\,pc.  7. The eastern jet is approaching, but is still partially covered along approximately 0.3\\,pc. This probably indicates a geometrically thick absorber, oriented roughly orthogonal to the jets.  8. Globally, the free-free opacity drops with increasing distance from the centre, but component variability suggests patchiness of the absorber. The highest turnover frequency of 43\\,GHz indicates a volume density of $n_{\\rm e}\\sim10^5$\\,cm$^{-3}$ if the free-free absorbing gas were distributed uniformly along a path-length of 0.5\\,pc.  9. H{\\sc i} atomic gas absorption is distributed in front of the approaching as well as the receding jet and has sub-pc scale structure. We distinguish three absorption systems with different characteristics.  10. ``High velocity'' H{\\sc i} gas, receding by 125--200\\,km\\,s$^{-1}$ with respect to the systemic velocity of NGC\\,1052 ($v_{\\rm sys}=1474$\\,km\\,s$^{-1}$), shows the most prominent absorption.  11. The high velocity H{\\sc i} absorber may have a continuous velocity gradient across the nucleus of some 10\\,km\\,s$^{-1}$\\,pc$^{-1}$, to connect a 20--25\\%\\ deep feature around 1620\\,km\\,s$^{-1}$ 1.5\\,pc away from the core along the receding jet, with 5--10\\%\\ deep features near 1640 and 1650\\,km\\,s$^{-1}$ located 1 and 1.5\\,pc away from the core along the approaching jet.  12. The innermost parsec (of the approaching jet) shows a deficit of this high-velocity component, which shows that it is local to the AGN. The atomic gas in the ``central hole\" may be largely ionised due to the proximity to the AGN; this is the location of the deepest free-free absorption.  13. A broad, shallow H{\\sc i} absorption component, with a peak depth of about 2\\%, spans the ``low velocity'' range 1440--1570\\,km\\,s$^{-1}$, asymmetrically straddling the systemic velocity. It covers all jet components in which we have the sensitivity to detect it, and could arise either in gas local to the AGN or on galactic scales.  14. Sharper H{\\sc i} features, having widths of 3--15\\,km\\,s$^{-1}$ and depths up to 5\\%, are superimposed in the range 1500--1570\\,km\\,s$^{-1}$. These features each extend over only a few tenths of a pc and can be detected at various places along the inner 2\\,pc of the approaching jet.  15. Since the absorbing neutral gas has structure on scales of 0.1--1\\,pc, back-lit by compact, near-relativistically moving features in the jets, it is logical that the overall H{\\sc i} absorption line profile of our data differs significantly from that in VLA observations taken 15 years earlier.  16. The 1667 and 1665\\,MHz 18\\,cm OH main lines are clearly present in absorption along the full velocity span seen also in H{\\sc i}, roughly 1440--1670\\,km\\,s$^{-1}$. The peak absorption depth in the 1667\\,MHz line is about 0.4\\%\\ at 1640\\,km\\,s$^{-1}$, suggesting a column depth of order $10^{14}$\\,cm$^{-2}$. The ratio between the OH 1667 and 1665\\,MHz lines ranges from near 1 at low velocities to approximately 2 in the high velocity system.  17. The 1612 and 1720\\,MHz 18\\,cm OH satellite lines have also been detected, in the high velocity system at 1630--1650\\,km\\,s$^{-1}$. The 1612\\,MHz line is in absorption, but the 1720\\,MHz line is in emission, both with a peak strength near 0.25\\%. Their conjugate profiles resemble the behaviour in \\object{Cen~A} and \\object{NGC\\,253}.  18. The OH main line and the total H{\\sc i} profiles are remarkably similar in the high velocity system, and we suggest co-location of these high velocity atomic and molecular gas components. They probably do not coincide with the H$_2$O masers at 0.1--0.2\\,pc along the receding jet, even though these are at the same velocity.  19. Other questions also remain regarding the nature of the high velocity system. Interpretation of the velocity gradient in H{\\sc i} as evidence for a structure rotating around the nucleus is contradicted by the fact that the centroid is redshifted by 150\\,km\\,s$^{-1}$ or more from the systemic velocity. But if it is instead infalling gas, the nature of the central hole in H{\\sc i} is unclear."
    },
    "0212/astro-ph0212340_arXiv.txt": {
        "abstract": "We investigate first order phase transitions from $\\beta$-equilibrated hadronic matter to color flavor locked quark matter in compact star interior. The hadronic phase including hyperons and Bose-Einstein condensate of $K^-$ mesons is described by the relativistic field theoretical model with density dependent meson-baryon couplings. The early appearance of hyperons and/or Bose-Einstein condensate of $K^-$ mesons delays the onset of phase transition to higher density. In the presence of hyperons and/or $K^-$ condensate, the overall equations of state become softer resulting in smaller maximum masses than the cases without hyperons and $K^-$ condensate. We find that the maximum mass neutron stars may contain a mixed phase core of hyperons, $K^-$ condensate and color superconducting quark matter. Depending on the parameter space, we also observe that there is a stable branch of superdense stars called the third family branch beyond the neutron star branch. Compact stars in the third family branch may contain pure color superconducting core and have radii smaller than those of  the neutron star branch. Our results are compared with the recent observations on RX J185635-3754 and the recently measured mass-radius relationship by X-ray Multi Mirror-Newton Observatory. ",
        "introduction": "The theoretical investigation of mass-radius relationship of compact stars is important because those could be directly compared with measured masses and radii from various observations. Consequently, the composition and equation of state (EoS) of dense matter in neutron star interior might be probed. Recently, a lot of interest has been generated about an isolated neutron star (INS) known as RX J185635-3754 \\cite{Pon}. Several interesting features of this INS have emerged from various observations by Chandra X-ray observatory and Hubble Space Telescope (HST). It is one of the closest isolated neutron stars to our Sun. It provides a great opportunity to measure its radius and put important constraints on the EoS of dense matter in its core because of its relative brightness, isolated nature and above all a thermal spectrum from X-ray to optical wave lengths. Already, one group who has analysed the Chandra data using a featureless blackbody spectrum, has claimed its radius is $\\sim 8$ km or less \\cite{Dra}. This implies a very soft EoS including exotic forms of matter. And their prediction is the compact star might be a quark star \\cite{Dra}. Alternatively, the Chandra data have been interpreted using a neutron star model where the gaseous outer layers of the neutron star turn into a solid due to a phase transition and this model gives an apparent radius $\\sim 10-12$ km \\cite{Zan}. On the other hand, the analysis of HST data on that compact star indicated a radius R=$11.4 \\pm  2$ km and a mass M=$1.7 \\pm 0.4 M_{solar}$ which may be explained by many EoS including hyperons and antikaon condensate \\cite{Lat02}. Recently three strong spectral lines in the spectra of 28 bursts of low mass X-ray binary EXO0748-676 have been observed by X-ray Multi Mirror-Newton Observatory. The estimated gravitational red shift from those lines is $z=0.35$ \\cite{Cot}. This provides us with many important informations about the structure of compact star because the gravitational redshift depends on the mass-radius ratio of a compact star.  The composition and structure of compact stars depend on the nature of strong interaction. Neutron star matter encompasses a wide range of densities- the density of iron nucleus at the surface to several times normal nuclear matter density in the core. Different phases of exotic matter with large strangeness fraction such as hyperon matter \\cite{Gle97}, Bose-Einstein condensates of strange mesons \\cite{Kap,Pra97,Sch96,Gle99,Pal,Bani00,Bani01,Bani02} and quark matter \\cite{Gle97,Pra97,Far84,Schr} may occur in neutron star interior. It was extensively investigated how the appearance of various forms of exotic matter influences the EoS and mass-radius relationship of compact stars. It is now well understood that each exotic component of dense matter makes the EoS soft. It was found in theoretical investigations \\cite{Lat01} that a soft EoS generally gave rise to a compact star with smaller maximum mass and radius than those of a stiffer EoS.  A transition from hadronic matter to deconfined strange quark matter is a possibility in neutron star interior. In earlier investigations, the quark phase was described by the MIT bag model \\cite{Gle97,Far84} in the hadron-quark phase transition. The recent development in dense matter physics points to the fact that the quark matter may be a color superconductor \\cite{Bar77,Frau,Bai,Alf98,Rap,Alf99,Ris}. In this case, quarks near their Fermi surfaces form Cooper pairs because quark-quark interaction is attractive in the antisymmetric color channel. The formation of diquark condensates breaks the color gauge symmetry. It was shown that at very high density quarks with all three flavors and  all three colors might pair up so as to produce an energetically favored state called the color-flavor-locked (CFL) phase\\cite{Raj01}. The color neutrality constraint is to be imposed in the CFL quark matter because a macroscopic chunk of quark matter must be color singlet \\cite{Raj02,Stei02}. Consequently, the CFL quark matter is charge neutral. Color and charge neutrality in the CFL quark matter require a non-zero value for the color chemical potential and electron chemical potential $\\mu_e=0$ respectively \\cite{Raj02,Stei02}. As the CFL condensate breaks chiral symmetry, the lightest degrees of freedom in this phase are Nambu-Goldstone bosons \\cite{Raj01,Red01,Bed02}. It was shown by various groups how the symmetric CFL phase behaves under stresses such as non-zero strange quark mass and electron chemical potential \\cite{Red01,Bed02}. The CFL phase could relax under those stresses forming meson condensation. One such possibility is $K^0$ condensation in the charge neutral CFL quark matter \\cite{Red01}.  Recently, nuclear-CFL quark matter phase transition \\cite{Alf01} and its impact on the structure of compact stars have been studied \\cite{Alf02}. Also, the structure of compact stars including pure CFL quark matter has been studied by others \\cite{Hov02}. Along with CFL quark matter, hyperons and $K^-$ condensate could exist in compact star interior. Many interesting things may happen if one form of exotic matter sets in before the other. The early appearance of hyperons was found to delay the hadron-unpaired quark matter phase transition or vice-versa \\cite{Pra97}. Similar effects of hyperons on the threshold of $K^-$ condensation in hadronic matter was observed by various groups \\cite{Pra97,Pal,Bani00,Bani01,Bani02}. Now the question is what could happen to nuclear-CFL quark matter phase transition if hyperons and $K^-$ condensate appear in compact star matter. So far no calculation of compact stars involving color superconducting quark matter, hyperons and $K^-$ condensate has been performed. In this paper, we investigate how the formation of hyperons and antikaon condensate in the hadronic matter influences the phase transition from hadronic matter to the CFL quark matter including $K^0$ condensate and the structure of compact stars. The paper is organised in the following way. In Sec. II, we describe the DDRH model for hadronic phase and also the CFL quark matter. Results of our calculation are explained in Sec. III. And Sec. IV provides a summary and conclusions. ",
        "conclusions": "We have investigated first order phase transitions from hadronic matter including hyperons and antikaon condensate to color-flavor-locked quark matter including the condensate of Goldstone boson $K^0$. The DDRH model has been adopted here to describe the hadronic phase. This model takes into account many body correlations in hadronic phase by density dependent meson-baryon couplings. Density dependent meson-baryon couplings are obtained from microscopic Dirac-Brueckner calculations using Groningen nucleon-nucleon potential. In this calculation meson-(anti)kaon couplings are density independent. Here $K^-$ condensation is found to be a second order phase transition. On the other hand, the CFL quark matter is described by the thermodynamic potential of Ref.\\cite{Raj01}. The role of strange quark mass and color-superconducting gap on the nuclear-CFL+$K^0$ matter phase transition has been studied for a fixed value of bag constant $B^ {1/4} =180$ MeV. The phase transition is delayed to higher density for a larger value of $m_s$. On the other hand, the phase transition sets in earlier as the value of gap increases. Also, we have found the early appearance of hyperons and /or Bose-Einstein condensate of $K^-$ mesons shifts the phase transition to higher density.  Equations of state for nuclear matter-CFL+$K^0$ matter phase transition have been studied here for different values of $m_s$ and $\\Delta$ and for a fixed value of bag constant. For a fixed value of gap, the EoS becomes stiffer in case of a larger value of $m_s$ resulting in larger maximum mass star. For a large value of gap such as $\\Delta = 100$ MeV, we find that different values of strange quark mass have no impact on the EoS and the maximum masses of compact stars. Similarly, a smaller value of gap softens the EoS leading to a smaller maximum mass star. We have also constructed equations of state including hyperons and $K^-$ condensate in the hadronic phase for different values of bag constant, gap, antikaon optical potential depth and strange quark mass. The EoS for NHQ, N$\\bar K$Q and NH$\\bar K$Q matter are softer than that of NQ matter. Consequently, the maximum star masses are smaller in the former cases. For all cases studied here, maximum mass neutron stars contain a mixed phase core of hadronic+CFL quark matter. There is a window in the parameter space for which all three exotic forms of dense matter i.e. hyperons, Bose-Einstein condensate of $K^-$ mesons and CFL quark matter are found to coexist in maximum mass neutron stars. In our calculation with CFL quark matter, maximum masses of compact stars range from 1.464 to 1.649 $M_{solar}$. Our results are consistent with the Hulse Taylor pulsar mass and recent observations. It is worth mentioning here that we have obtained stable branches of superdense stars called third family branches beyond neutron star branches for certain combinations of parameters. Superdense stars in the third family branch contain a pure CFL quark matter core. The compact stars in the third family have smaller radii and different compositions than those of the neutron star branch.  In this calculation, we do not consider the density dependence of strange quark mass, gap and bag parameter. Therefore, it would be interesting to investigate how density dependent strange quark mass, gap and bag parameter in the Nambu-Jona-Lasinio model \\cite{Stei02} modify our results.  \\centerline{\\bf Acknowledgements} D.B. is grateful to Prof. (Dr.) W. Greiner for the warm hospitality of the Institut f\\\"ur Theoretische Physik, Frankfurt Universit\\\"at where a part of the work was completed and the Alexander von Humboldt Foundation for the support.  \\newpage"
    },
    "0212/astro-ph0212389_arXiv.txt": {
        "abstract": "We investigate the nature of the faint X-ray source population through X-ray spectroscopy and variability analyses of 136 AGN detected in the 2 Ms Chandra Deep Field-North survey with $> 200$ background-subtracted 0.5--8.0~keV counts [$F_{\\rm 0.5-8.0~keV}=(1.4$--$200)\\times10^{-15}$~erg~cm$^{-2}$~s$^{-1}$]. Our preliminary spectral analyses yield median spectral parameters of $\\Gamma=1.61$ and intrinsic $N_{\\rm H}=6.2\\times10^{21}$~cm$^{-2}$ ($z=1$ assumed when no redshift available) when the AGN spectra are fitted with a simple absorbed power-law model. However, considerable spectral complexity is apparent (e.g., reflection, partial covering) and must be taken into account to model the data accurately. Moreover, the choice of spectral model (i.e., free vs. fixed photon index) has a pronounced effect on the derived $N_{\\rm H}$ distribution and, to a lesser extent, the X-ray luminosity distribution. Ten of the 136 AGN ($\\approx7$\\%) show significant Fe K$\\alpha$ emission-line features with equivalent widths in the range 0.1--1.3~keV. Two of these emission-line AGN could potentially be Compton thick (i.e., $\\Gamma < 1.0$ and large Fe K$\\alpha$ equivalent width). Finally, we find that 81 ($\\approx60$\\%) of the 136 AGN show signs of variability, and that this fraction increases significantly ($\\approx80$--90\\%) when better photon statistics are available. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/astro-ph0212495_arXiv.txt": {
        "abstract": "We have observed the cosmic microwave background (CMB) in three regions of sky using the Very Small Array (VSA) in an extended configuration with antennas of beamwidth $2^{\\circ}$ at 34~GHz. Combined with data from previous VSA observations using a more compact array with larger beamwidth, we measure the power spectrum of the primordial CMB anisotropies between angular multipoles $\\ell = 160$ -- $1400$. Such measurements at high $\\ell$ are vital for breaking degeneracies in parameter estimation from the CMB power spectrum and other cosmological data. The power spectrum clearly resolves the first three acoustic peaks, shows the expected fall off in power at high $\\ell$ and starts to constrain the position and height of a fourth peak. ",
        "introduction": "Acoustic peaks in the power spectrum of cosmic microwave background~(CMB) anisotropies were predicted by \\citet{sakharov}, \\citet{sunyaev-70} and \\citet{peebles-70}. Recently several experiments have accurately measured the first of these peaks and have detected the second peak ~\\citep{lee-01,netterfield-01,halverson-02,VSApaperIII}.  Importantly, despite the differing types of potential systematic errors suffered by these experiments, they are in very good agreement with each other. The CBI experiment~\\citep{padin-02} has measured the power spectrum out to $\\ell \\ \\sim \\ 3500$~\\citep{mason-02,pearson-02}, and has detected the predicted fall in power level due to incoherent addition of temperature fluctuations along the line of sight and Silk damping~\\citep{silk}, but does not have the $\\ell$-resolution to define the peak structure. Resolving the third and subsequent CMB peaks is essential to constrain further cosmological models.  The Very Small Array (VSA)~\\citep{VSApaperI} (hereafter Paper\\,I) is an interferometric array which has measured the CMB power spectrum between $ \\ell \\ = $ 150--800 as described in \\citet{VSApaperII} and \\citet{ VSApaperIII} (hereafter Paper\\,II and Paper\\,III respectively).  These measurements were made with the VSA's 14 antennas arranged in a compact configuration. The VSA has since been upgraded with larger apertures to allow observations in an extended configuration giving high temperature-sensitivity measurements on smaller angular scales. This paper describes the first VSA measurements with the extended array and the power spectrum from the combined compact and extended arrays.  \\begin{table} \\begin{center} \\caption{VSA field positions and total effective integration time remaining after flagging and filtering of the data.\\label{tab:fieldpos}} \\begin{tabular}{lccc} \\hline & RA (J2000) & DEC (J2000)& $\\rm T_{int}$ (hrs)\\\\ \\hline VSA1E     &     00 22 37  & 30 16 38 & 106\\\\ VSA1F     &     00 16 52  & 30 24 10 & 94\\\\ VSA1G     &     00 19 22  & 29 16 39 & 79\\\\ VSA2E     &     09 37 57  & 30 41 28 & 110\\\\ VSA2F     &     09 43 46  & 30 41 14 & 101\\\\ VSA2G     &     09 40 53  & 31 46 21 & 115\\\\ VSA3E     &     15 31 43  & 43 49 53 & 130\\\\ VSA3F     &     15 38 38  & 43 50 18 & 114\\\\ VSA3G     &     15 35 13  & 42 45 05 & 112\\\\ \\hline \\end{tabular} \\end{center} \\end{table} ",
        "conclusions": "We have measured the power spectrum of the CMB over the multipole range $\\ell = 160$ -- $1400$ using two configurations of the VSA. The combined power spectrum clearly shows three peaks, a sharp fall-off in power above $\\ell~\\sim~800$, and marginal evidence for a fourth peak. Figure~\\ref{CMB_comparison} shows the VSA data plotted alongside that from other recent CMB experiments: there is excellent agreement between all the experiments.  To quantify the significance of these peaks in the power spectrum, we used a model of five Gaussians parametrised by their amplitude, position and width, as described in \\citet{Odman2002}. We analysed the VSA data alone, plus two data sets consisting of a compilation of recent data, and used an MCMC routine to constrain the parameters simultaneously. In Table~\\ref{constraints}, we list the constraints on the amplitudes and positions of the first four peaks after marginalising over their widths (there are no interesting constraints on the fifth peak). Including the new VSA measurement improves the constraints on the amplitude of the fourth peak, as well as confirming the detection of the first three peaks.  Fits of adiabatic inflationary model power spectrum, with associated estimates of cosmological parameters, are presented in a companion paper \\citep{VSApaperVI}.  In our current observing configuration we are increasing the size of the mosaiced regions and also observing new regions; this will increase both the sensitivity and $\\ell$-resolution of the power spectrum. We also plan to increase the $\\ell$-range of the VSA using larger antennas and longer baselines to study further peaks in the primordial spectrum and give information on the region where CBI observes excess power~\\citep{mason-02}.  Band powers, correlation matrices and window functions can be downloaded from\\\\ \\noindent \\verb+http://mrao.cam.ac.uk/telescopes/vsa/results.html+  \\begin{table} \\begin{center} \\caption{Constraints on parameters of a multiple Gaussian fit to the power spectrum data. The inclusion of the new VSA data places constraints on the power beyond the first three peaks. Dataset I consists of the QMASK compilation~\\citep{Xu2002}, MSAM~\\citep{Wilson2000}, Python V~\\citep{Coble2001}, Boomerang 98~\\citep{netterfield-01}, Maxima~\\citep{lee-01}, DASI~\\citep{halverson-02}, CBI~\\citep{pearson-02}, Archeops~\\citep{Benoit2002} and the VSA data from the compact array configuration. Dataset II is identical, except for the inclusion of the new VSA data. Note that, while the lack of a lower bound on a $\\Delta$T implies a significant contribution to the power from adjacent Gaussians, narrow constraints on peak locations imply strong evidence for spectral features rather than a flat spectrum of excess power.\\label{constraints}} \\begin{tabular}{|l|c|c|c|}\\hline & VSA only & dataset I & dataset II \\\\ \\hline $\\Delta T_1$ & $71.4_{-5.5}^{+11.3}$  & $70.2_{-2.0}^{+2.1}$  & $71.7_{-2.1}^{+2.0}$\\\\ $\\Delta T_2$ & $45.5_{-5.6}^{+6.4}$  & $38.5_{-7.8}^{+8.1}$  & $42.3_{-10.4}^{+4.7}$\\\\ $\\Delta T_3$ & $47.0_{\\mbox{\\tiny unbounded}}^{+8.9}$  & $42.5_{-2.2}^{+1.5}$ & $43.3_{-2.3}^{+1.7}$ \\\\ $\\Delta T_4$ & $25.9_{\\mbox{\\tiny unbounded}}^{+14.1}$  & $18.0_{\\mbox{\\tiny unbounded}}^{+19.1}$ & $19.6_{-11}^{+11.7}$ \\\\ $\\ell_1$ & $224_{-28}^{+11}$  & $214_{-5}^{+4}$  & $216_{-6}^{+3}$\\\\ $\\ell_2$ & $509_{-84}^{+81}$  & $511_{-16}^{+12}$  & $510_{-14}^{+11}$\\\\ $\\ell_3$ & $771_{-70}^{+55}$  & $786_{-48}^{+29}$  & $782_{-52}^{+23}$\\\\ $\\ell_4$ & $1171_{-205}^{+67}$  & $1260_{-321}^{+37}$  & $1157_{-118}^{+167}$\\\\ \\hline \\end{tabular} \\end{center} \\end{table}  \\begin{figure*} \\centering \\begin{minipage}{1\\textwidth} \\centering \\epsfig{file=final_extended_bw.ps,height=170mm,angle=270,clip=} \\caption{Combined CMB power spectrum from the three mosaiced VSA fields.  The error-bars represent $1\\sigma$ limits; the two sets of data points correspond to alternative interleaved binnings of the data. \\label{CMB_doublebin}} \\end{minipage}\\\\[20pt] \\begin{minipage}{1\\textwidth} \\centering \\epsfig{file=final_combined_bw.ps,height=170mm,angle=270,clip=} \\caption{A comparison of the VSA data (filled circles) with results from the BOOMERANG (open squares), MAXIMA (open circles), DASI (open triangles) and CBI (stars) experiments. Two sets of error bars are plotted for each data set; the smaller (inner) of the two indicate only random errors, whilst the larger bars indicate the amount by which the inner points could move due to absolute calibration and beam uncertainty. In each case the error bars indicate $1\\sigma$ limits.\\label{CMB_comparison}} \\end{minipage} \\end{figure*}"
    },
    "0212/astro-ph0212176_arXiv.txt": {
        "abstract": "{ Five years of \\eros\\ data towards the Small Magellanic Cloud have been searched for gravitational microlensing events, using a new, more accurate method to assess the impact of stellar blending on the efficiency.  Four long-duration candidates have been found which, if they are microlensing events, hint at a non-halo population of lenses. Combined with results from other \\eros\\ observation programs, this analysis yields strong limits on the amount of Galactic dark matter made of compact objects. Less than 25\\% of a standard halo can be composed of objects with a mass between $2\\times10^{-7}\\,\\rm M_\\odot$ and 1 $\\rm M_\\odot$ at the 95\\% C.L.  \\keywords {Galaxy: halo -- Galaxy: kinematics and dynamics -- Galaxy: stellar content -- Magellanic Clouds -- dark matter -- gravitational lensing } }  \\authorrunning{EROS Collaboration} \\titlerunning{Limits on Galactic Dark Matter with 5 Years of EROS SMC Data}  \\def\\ie{{\\em i.e.}} \\def\\eros{{\\sc eros}} \\def\\macho{{\\sc macho}} \\def\\lmc{{\\sc lmc}} \\def\\smc{{\\sc smc}} \\def\\sm98{{\\sc sm{\\footnotesize 98}}} \\def\\ccd{{\\sc ccd}} \\def\\etal{{et al.}} ",
        "introduction": " ",
        "conclusions": "The analysis of five years of microlensing data towards the \\smc\\ yielded four candidates, and allowed to put stringent limits on the amount of galactic dark matter made of compact objects. Objects with a mass between $2\\times10^{-7}\\,\\rm M_\\odot$ and 1 $\\rm M_\\odot$ cannot account for more than 25\\% of the mass of a standard spherical, isothermal and isotropic Galactic halo of $4 \\times 10^{11} \\,\\rm M_\\odot$ out to 50~kpc.  The new method described in this paper to combine the results and use all the information available on the events has one drawback compared to methods used in previous works: it could be more sensitive to the halo model, adding an extra dependence on the velocity distribution. Previously, only the influence of the halo model on the number of expected events has been studied, as in \\sm98. A change in the velocity distribution, however, would mainly shift the ``dent'' in the \\eros2 \\lmc\\ limit. In addition, since the final exclusion limit is quite flat over a large range of masses, the final result should not be influenced by reasonable changes in the velocity distribution assumptions.   This limit excludes a sizeable fraction of the allowed domain shown in Fig.~12 of Alcock et al. (2000).  All the candidates that have been detected so far towards the \\smc\\ have long durations, and seem more compatible with unidentified variable stars or self-lensing within the cloud than with halo objects. These would need to have supersolar masses to account for such durations.  The statistics are still very low for a quantitative comparison of the \\smc\\ and \\lmc\\ samples of microlensing  candidates. They have, however, quite different characteristics which is poorly compatible with a unique population of lenses."
    },
    "0212/astro-ph0212162_arXiv.txt": {
        "abstract": "{We make use of two empirical relations between the black hole mass and the global properties (bulge luminosity and stellar velocity dispersion) of nearby elliptical galaxies, to infer the mass of the central black hole ($\\CMcal{M}_{BH}$) in low redshift radiogalaxies.  Using the most recent determinations of black hole masses for inactive early type galaxies we show that the bulge luminosity and the central velocity dispersion are almost equally correlated (similar scatter) with the central black-hole mass. Applying these relations to two large and homogeneous datasets of radiogalaxies we find that they host black-holes whose mass ranges between $\\sim 5\\times10^7$ to $\\sim 6\\times10^9\\CMcal {M}_{\\odot}$ (average $<Log\\CMcal{M}_{BH}> \\sim$8.9). $\\CMcal{M}_{BH}$ is found to be proportional to the mass of the bulge ($\\CMcal{M}_{bulge}$).  The distribution of the ratio $\\CMcal{M}_{BH}$/$\\CMcal{M}_{bulge}$ has a mean value of 8$\\times10^{-4}$ and shows a scatter that is consistent with that expected from the associated errors. At variance with previous claims no significant correlation is instead found between $\\CMcal{M}_{BH}$ (or $\\CMcal{M}_{bulge}$) and the radio power at 5 GHz. ",
        "introduction": "There is a large consensus about the existence of super massive black holes (SBHs) at the center of nearby inactive galaxies as well as in the nuclei of active galaxies and quasars (see e.g. for a recent review \\cite{ferrarese}).  A large body of data, in particular based on high resolution HST observations, is now available (see e.g. \\cite{kormendy2}) to support the presence of such massive black holes (BH) using different techniques.  It is believed that SBHs play an important role in the formation and evolution of massive galaxies and also to be a key component for the development of the nuclear activity. However in spite of this apparently ubiquitous presence of SBHs in galaxies, our understanding on how the galaxies and their central BHs are linked in the process of formation of the observed structures is still poor (see \\cite{silk}; \\cite{haehnelt}; \\cite{adams}).  The most important result obtained from the measured BH masses in nearby galaxies is the existence of a significant correlation between the black hole mass ($\\CMcal{M}_{BH}$) and the mass of the bulge component ($\\CMcal{M}_{bulge}$) of the host galaxy ($\\CMcal{M}_{bulge}$). From the observational point of view this correlation is translated into relationships between $\\CMcal{M}_{BH}$ and bulge luminosity $L_{bulge}$ (\\cite{magorrian}, \\cite{kormendy2}) and between $\\CMcal{M}_{BH}$ and the stellar velocity dispersion $\\sigma$ (\\cite{FM}; \\cite{gebhardt} ).  Although based on a small number ($\\sim$30) of nearby galaxies for which direct dynamical measurements of $\\CMcal M_{BH}$ have been secured, and in spite of their scatter [$\\sim$0.4 in $Log (\\CMcal{M}_{BH})$], these empirical relationships offer a new tool for evaluating $\\CMcal{M}_{BH}$ in various types of AGN, provided that bulge luminosities and/or velocity dispersion be available (see also \\cite{MD}, \\cite{FCT}).  In this paper we make use of such relationships to investigate the BH mass distribution of two large and homogeneous datasets of low redshift radiogalaxies (RG) for which we have previously studied the morphological, structural, photometrical and kinematical properties (\\cite{FFS}; ~\\cite{fede1},b; ~\\cite{bettoni}).  The derived BH masses of radiogalaxies are then used to investigate the connections between $\\CMcal{M}_{BH}$, the mass of the galaxy and the radio power.  To this aim we first describe our samples (Section 2) and revisit the relations $\\CMcal{M}_{BH}$-$L_{bulge}$ and $\\CMcal{M}_{BH}$-$\\sigma$ for nearby early-type galaxies (Section 3). Then we use these relationships to evaluate $\\CMcal{M}_{BH}$ of radio galaxies (Section 4) and to study the connections between $\\CMcal{M}_{BH}$ and the mass of the bulge component of the host galaxy and between $\\CMcal{M}_{BH}$ and the radio luminosity. A summary of the main conclusions of this study is reported in Section 5. In our analysis we assume H$_{0}$=50 Km s$^{-1}$ Mpc$^{-1}$ and $\\Omega_{0}$=0. ",
        "conclusions": "We used of empirical relations between $\\CMcal M_{BH}$ and either the velocity dispersion ($\\sigma$) or the bulge luminosity M$_R$(bulge) derived for nearby early type galaxies to infer the mass of the central BH of low redshift radio galaxies.  The main conclusions of this study are:  1) Using only nearby galaxies of E-type morphology, for which BH masses are available, the two relationships $\\CMcal M_{BH}$ - $\\sigma$ and $\\CMcal M_{BH}$ - M$_{bulge}$ exhibit a similar scatter ($\\sim$0.4 in Log$\\CMcal M_{BH}$) which is also consistent with that expected from the estimated errors  on the involved parameters.  2) We showed that when the dependence on the adopted cosmology (for the $\\CMcal M_{BH}$ - M$_{bulge}$ relation) is properly taken into account the two relations predict the same value of the central BH mass (within the expected uncertainties). This means that the BH mass can be reliably estimated from the observable parameter M$_R$(bulge) which is more easily measurable than $\\sigma$.  3) The inferred BH mass of low redshift radio galaxies is in the range 5$\\times10^7$ to 5$\\times10^9$ $\\CMcal {M}_{\\odot}$.  4) We found that the central BH mass is linearly correlated with that of its host galaxy the average Black Hole to bulge mass ratio is $<Log[\\CMcal{M}_{BH}/\\CMcal{M}_{bulge}]>\\sim$ -3.1.  5) The total (or core) radio power at 5~GHz is not correlated neither with the mass of the central BH nor with that of the galaxy; the radio power is always in excess by 2-3 order of magnitude with respect to what would be expected from low accretion rate models (ADAF)."
    },
    "0212/astro-ph0212448_arXiv.txt": {
        "abstract": "We discuss the production of the element $^6$Li in the early Galaxy by cosmic rays accelerated at structure formation shocks, driven by the hierarchical merging of sub-Galactic halos during Galaxy formation. The salient features of this scenario are discussed and compared with observations of $^6$Li in metal-poor halo stars, including a recent Subaru HDS result on the star HD140283. Some unique predictions of the model are clearly testable by future observations and may also provide important insight into how the Galaxy formed. ",
        "introduction": "The origin of the light element isotope $^6$Li in old, metal-poor halo stars (MPHS) is currently mysterious (see \\cite{vca00,hob00,rlk00,nis00} for reviews). Although the most widely discussed models of light element production based on nuclear reactions by cosmic rays (CRs) from supernovae (SNe) give a good account of the Be and B observed in such stars, they generally fall short for the observed $^6$Li. Such models must resort to rather implausible or contrived assumptions for $^6$Li, e.g. a CR injection efficiency substantially higher than normally inferred \\cite{rslk00,sy01} or an additional low energy CR component lacking observational support \\cite{van99}. We have recently proposed a fundamentally different $^6$Li production scenario: nuclear reactions induced by CRs accelerated at structure formation (SF) shocks, i.e. gravitational virialization shocks driven by the infall and merging of sub-Galactic halos during the hierarchical build-up of structure in the early Galaxy \\cite{si02}. Given below is a very brief description of the scenario, along with a discussion in relation to recent and future observations. See \\cite{si02,aok02} for more details. ",
        "conclusions": ""
    },
    "0212/astro-ph0212354_arXiv.txt": {
        "abstract": "N66/NGC\\,346 is the largest and brightest \\HII\\ region in the Small Magellanic Cloud and contains at least one known supernova remnant \\SNR.  Optical emission from the remnant is overwhelmed by the bright photoionized emission from the nebula, but the remnant has been detected by way of far ultraviolet absorption lines.  Here we present data from the Far Ultraviolet Spectroscopic Explorer ({\\it FUSE}) satellite showing strong \\OVI\\ and \\CIII\\ emission from a position at the edge of \\SNR.  We also present high-resolution, long-slit \\Ha\\ spectra across N66 showing high- and low-velocity emission corresponding closely to the X-ray boundaries of the supernova remnant.  We use these FUV and optical data to determine the physical parameters of the shock and interaction geometry with N66.  We find that ionizing photons from the many massive cluster stars nearby likely affect the ionization balance in the post-shock gas, hindering the production of lower-ionization and neutral species.  We discuss the importance and inherent difficulty of searching for supernova remnants in or near bright \\HII\\ regions and suggest that the far ultraviolet provides a viable means to discover and study such remnants. ",
        "introduction": "The giant \\HII\\ region N66 is the largest, brightest star forming region in the Small Magellanic Cloud (SMC).  The populous star cluster NGC\\,346 embedded in the nebula is home to several dozen early type stars including a dozen of spectral types O7 and earlier.  Two arcminutes to the east lies the massive, luminous WN+OB system HD\\,5980 which underwent a luminous blue variable  (LBV) outburst in 1994 \\citep{Koenigsberger95,Barba95}.  These stars are the main source of ionization in the nebula.  N66 also contains at least one known supernova remnant (SNR), \\SNR.  {\\it Einstein} X-ray observations showed a bright source in N66 tentatively identified as a SNR \\citep{IKT,WangWu92}.  Radio observations by \\citet{Mills82} showed bright, extended, thermal emission at 408, 843 and 5000 MHz corresponding to the \\HII\\ region, but no clear sign of the SNR.  Optical observations also showed only the extremely bright nebular emission but no filamentary structure associated with the SNR.  However, \\citet{ChuKennicutt88} discovered faint, high-velocity \\Ha\\ emission in a long-slit spectrum across N66 and suggested the existence of a SNR.  Subsequently, \\citet{YKT} subtracted smoothed \\Ha\\ data from 843 MHz radio continuum observations and revealed a remnant $\\sim$3.2\\arcmin\\ ($\\sim$55 pc at an SMC distance of 59 kpc) in diameter and confirmed the remnant \\SNR.  X-ray spectroscopy of \\SNR\\ with ASCA \\citep{Yokogawa00} showed that the spectrum can be modeled as a hot thermal plasma.  High-resolution X-ray observations of N66 with the {\\it Chandra X-ray Observatory} \\citep{Xmega} show a 130\\arcsec$\\times$100\\arcsec\\ region of extended, center-bright, thermal emission corresponding to the SNR.  These high-resolution data show \\SNR\\ to have relatively uniform surface brightness and no clear temperature gradient from center to rim.  Difference imaging techniques applied to recent radio observations \\citep{Filipovic03} at 1.42, 2.37, and 4.80 GHz (20, 13 and 6 cm) show \\SNR\\ to be a non-thermal, limb-brightened shell with a spectral index of approximately $\\alpha$=$-$0.17 (S$_\\nu\\propto\\nu^\\alpha$).  The X-ray and radio emitting regions are similar in extent at a resolution of $\\sim$20\\arcsec.  Hence, the \\citet{YKT} estimate of angular size is probably too large, and the X-ray/radio SNR is roughly 36 $\\times$ 28 pc in extent. Thus, \\SNR\\ may be a member of the ``mixed morphology'' class of SNRs \\citep{RhoPetre98}, with a radio shell and filled-center X-ray emission, although the presence of HD~5980 along the sight line confuses the situation. \\citet{Xmega} resolve X-ray emission from HD~5980 itself for the first time, and note similarities with the galactic object $\\eta$ Carinae, which has diffuse surrounding X-ray emission \\citep{Sew81}.  \\citet{Xmega} mention a possible association of HD~5980 and \\SNR, but FUV absorption studies have clearly determined that HD~5980 is behind the SNR.  FUV observations toward HD\\,5980 show the expected Galactic and SMC absorption at $\\rm v_{lsr}\\approx$0 \\kms\\ and $\\rm v_{lsr}\\approx$$+$150 \\kms, respectively, as well as high-velocity absorption indicating the receding shell of an intervening SNR.  Absorption at $+$300 \\kms\\ was found in IUE spectra by \\citet{deBoerSavage80}; this was later confirmed by \\citet{FitzpatrickSavage83} who suggested an intervening SNR.  \\citet{Koenigsberger01} detected absorption systems at $+$300 and $+$330 \\kms\\ in \\CIV, \\NV\\ and several other ions in high-resolution {\\it HST} STIS spectra.  They also see very weak components at $+$20 and $+$50 \\kms, possibly consistent with the approaching side of the remnant if indeed these features are real.   \\citet{Hoopes01} also see absorption at $\\sim$$+$300 \\kms\\ in \\CIII\\ and \\OVI\\ in {\\it FUSE} observations of HD\\,5980 and conclude that the star must lie behind the receding shell of \\SNR.  \\citet{Hoopes01} also find that a {\\it FUSE} spectrum of Sk80 (AV232), an O7Ia star $\\sim$1\\arcmin\\ southeast of HD\\,5980 but also seen in projection within the X-ray emission, shows none of the high-velocity absorption seen toward HD\\,5980.  Since Sk80 lies closer to the X-ray edge of the remnant, we might expect a lower line-of-sight velocity for SNR gas in this sight line.  However, the \\OVI\\ absorption profile near the SMC systemic velocity ($\\sim$150 \\kms) is indistinguishable from that of HD\\,5980.  Thus, Sk80 must either lie within or in front of the remnant.  Far-ultraviolet (FUV) observations are useful in the study of SNRs.  Several strong emission lines including \\CIV\\ $\\lambda\\lambda$1548,1550, \\NV $\\lambda\\lambda$1238,1242, and especially \\OVI\\ $\\lambda\\lambda$1032,1038 are good shock diagnostics and are produced in regions where the temperature is lower than typical X-ray producing regions and higher than regions of bright optical emission \\citep{SutherlandDopita93}.  \\OVI\\ emission, in particular, arises in gas at temperatures near $3\\times10^5$ K--a condition almost never reached by photoionization--and is thus a sensitive tracer of shock-heated material.  In this paper we present a {\\it FUSE} observation of emission at the X-ray edge of \\SNR\\  that shows strong emission in \\OVI\\ $\\lambda\\lambda$1032,1038 and \\CIII\\ $\\lambda$977.  We also present new longslit echelle observations showing high- and low-velocity material in \\Ha\\ across the face of \\SNR.  The observations are described in \\S2.  In \\S3, we discuss the SNR kinematics as revealed by the optical and FUV data.  In \\S4, we compare the \\OVI, \\CIII, and \\Ha\\ measurements with the predictions of shock models to derive shock velocity, preshock density and ram pressure at one location in the SNR.  Using the same models with the FUV absorption data of \\citet{Hoopes01} and \\citet{Koenigsberger01}, we derive similar values for material on the sight line toward HD\\,5980.  We propose a physical picture for \\SNR\\ in relation to N66.  Finally we discuss the implications these observations have for the detection of SNRs in bright nebulae and OB associations.  A summary is presented in \\S5. ",
        "conclusions": "%  The detection of \\OVI\\ emission as well as the kinematics present in the \\Ha\\ echelle data imply the presence of a strong shock in N66 associated with \\SNR.  The SNR shock model was also favored by \\citet{Hoopes01} for the high velocity absorption towards HD\\,5980.  Detection of the front side of \\SNR\\ has not been clearly established because the low velocity components cited by \\citet{Koenigsberger01} were detected at such a marginal level. In any event, the near side of the SNR is considerably more difficult to detect both in emission and absorption than the back side.  Optical emission from the \\HII\\ region complicates the situation and masks the morphology of the SNR.  In the following discussion we adopt a working model of the physical association of various components in the N66 region.  A roughly spherical SNR is located on the near side of N66.  The rear side of the SNR is propagating into relatively denser nebula while the near side is propagating through the more rarefied SMC ISM.  HD\\,5980 lies behind the remnant, embedded within N66.  The NGC\\,346 cluster stars lie outside the SNR, also embedded within N66.  We determine the shock velocity, preshock hydrogen density and ram pressure of the shock wave using the emission and absorption sight lines toward \\SNR.  By examining the surface brightness ratio between \\OVI\\ and \\Ha, we measure the ``completeness'' of the shock.  Then we discuss the relationship between the SNR and its surroundings and the implications of this analysis for observing SNRs in other \\HII\\ regions.  \\subsection{Physical Properties}  %  In the N66 reference frame (v$\\rm_{lsr}$=$+$158 \\kms), the material at Position~1 has a radial velocity of $+$70 \\kms\\ while toward HD\\,5980, $\\rm v_r=+142$ \\kms.  From X-ray and radio images, we choose the center of the remnant to be (J2000) $\\alpha$=00:59:27, $\\delta$=$-$72:10:15, about 20\\arcsec\\ south of HD\\,5980.  This position is also consistent with the location of the highest-velocity \\Ha\\ emission (Figure~3).  Using the projected distances from this center and the two radial velocity measurements, we derive a radius of curvature for the shock front of 80$^{+19}_{-10}$\\arcsec\\ and an expansion velocity for the \\OVI\\ gas of v$\\rm_{exp}=147\\pm7$ \\kms.  This places the observed X-ray edge on the eastern side of the shock front inclined $\\sim$60\\dd\\ to our line of sight.  The expansion velocity of 147 \\kms\\ is not the shock speed, but rather the bulk velocity of the postshock material which should be between 0.75 and 1.0 times the shock speed.  First we caculate shock models for comparison against the observations.  The models were calculated using an updated version of the code described by \\citet{Raymond79}.  The main input parameters are the shock velocity, the preshock density and the elemental abundances.  We ran models for shock velocities between 140 and 200 \\kms\\ spaced by 10 \\kms.  A preshock hydrogen number density of 1 cm$^{-3}$, and SMC elemental abundances \\citep{RussellDopita90} were used in all the models.  The calculation is followed until the the recombination zone is complete and the gas temperature has reached about 1000 K.  We present the intensities of selected lines from the shock models in Table~3.  The intensities are relative to \\Ha=100 and the flux from nearby multiplet lines of an ion have been summed.  The \\OVI\\ line flux produced by a shock increases sharply as a function of shock velocity in the range 150 -- 200 \\kms.  Shocks below $\\sim$160 \\kms\\ do not produce much \\OVI\\ and thus this represents a lower limit to the actual shock velocity in the SNR.  Over this same range of shock velocities, the \\CIII\\ flux remains nearly constant.  The changing ratio of I(\\OVI)/I(\\CIII) is due to the increasing ionization of oxygen at higher velocities.  As discussed in \\S3.1, the \\CIII\\ is strongly affected by absorption and scattering, and so it is not possible for us to derive an accurate assessment of the intrinsic ratio of I(\\OVI)/I(\\CIII) in the spectrum of Position~1.  However, we note that using the total (corrected) \\OVI\\ flux with the observed \\CIII\\ flux places an upper limit of F(\\OVI)/F(\\CIII)$\\le$3.  Ignoring the effect of differential reddening for the moment, this clearly points us toward the lower end of the velocity range where \\OVI\\ is produced.  Based on this ratio, the observed expansion velocity and the shock velocity requirements for \\OVI\\ production, we adopt 160 \\kms\\ as a representative shock velocity at Position~1 and use details from the model at this velocity below.  We now compare the observed and calculated \\OVI\\ surface brightnesses to estimate the preshock density required to produce the observed \\OVI\\ flux.  The model predicts an \\OVI\\ surface brightness (in both lines of the doublet) of I$\\rm_{mod}=6.08\\times10^{-5}$ \\flux\\ into 2$\\pi$ steradians (for v$\\rm_{shock}$=160 \\kms, $\\rm n_o$=1 cm$^{-3}$, He/H=0.08) or I$\\rm_{mod}=2.28\\times10^{-16}$ \\surfb.  In the optically thin limit, intensity will scale with density so we can write  \\[ \\frac{I_{obs}~C_{1035}}{A}~=I_{mod}~n_o\\]  \\noindent where $C_{1035}$ is the reddening correction factor at 1035\\AA, $A$ is the surface area of shock (in arcseconds$^2$) within the aperture and $n_o$ is the preshock density in units of cm$^{-3}$.  I$\\rm_{obs}$ is the total \\OVI\\ flux in both lines which we estimate by multiplying the corrected \\OVI\\ $\\lambda$1032 flux by 1.5, thus I$\\rm_{obs}=1.91\\times10^{-13}$ \\flux.  From Figure~1, we estimate that $\\sim$$\\frac{1}{3}$ of the LWRS aperture is filled with X-ray emission.  If the \\OVI\\ emission follows the same pattern and the shock front is tilted 60\\dd\\ from our line of sight, then the aperture sampled emission from 600 arcseconds$^2$ of shock front.  Reddening is a potential source of uncertainty.  \\citet{MPG} found the mean E(B$-$V)=0.14 for the NGC\\,346 stellar population.  However, since the SNR lies in front of the bulk of the cluster, this may overestimate the extinction to the SNR.  On the other hand, using the spectral type and observed colors of Sk80, one can derive E(B$-$V)=0.11.  We adopt this value and assume R$\\rm_{v}$=3.1 (typical for the diffuse ISM) and the extinction curve of \\citet{HutchingsGiasson01} (E(1035\\AA$-$V)/E(B$-$V)=18).  This yields a reddening correction C$_{1035}$=4.5.  Solving the above equation, we get n$\\rm_o=6.3$ cm$^{-3}$ at Position~1, similar to an azimuthally-averaged electron density of 7 cm$^{-3}$ at the radius of Position~1 derived by \\citet{Kennicutt84}.  The dynamical pressure is then $\\rho~v_s^2=1.3~m_H~n_o~v_s^2=3.5\\times10^{-9}$ dyne cm$^{-2}$.  If more reddening is assumed and E(B$-$V) is raised to 0.14 (C$_{1035}$=6.8), the changes are minor:  n$\\rm_o=9.5$ cm$^{-3}$ and $\\rm P=5.2\\times10^{-9}$ dyne cm$^{-2}$.  Less reddening (E(B$-$V)=0.08, C$_{1035}$=3.0) would yield n$\\rm_o=4.2$ cm$^{-3}$, $\\rm P=2.3\\times10^{-9}$ dyne cm$^{-2}$.  This pressure is higher than those seen in regions of the Vela SNR.  \\citet{Sankrit01} derive n$\\rm_o=0.5$ cm$^{-3}$ and P=3.7$\\times10^{-10}$ dyne cm$^{-2}$ for a X-ray bright, nearly face-on shock.  \\citet{Raymond97} find P=1.6$\\times10^{-9}$ dyne cm$^{-2}$ for an edge-on shock.  \\citet{JenkinsWallerstein95} find P=2-4$\\times10^{-10}$ dyne cm$^{-2}$ for another locations in Vela while \\citet{Kahn85} find P$\\sim$1$\\times10^{-10}$ dyne cm$^{-2}$ for a region of diffuse X-ray emission.  This variation in pressures within the same remnant suggests local density and pressure enhancements caused by reverse shocks from blastwave-cloud interactions.  We have only two widely-separated measurements for \\SNR\\ and can say nothing so specific but the high pressure and density at Position~1 as well as the X-ray emission morphology suggest similar shock-cloud interactions.  The shock models discussed above predict the total column depths of ions produced in the post-shock flow.  \\citet{Hoopes01} measured the column depths of several highly ionized species in the high-velocity component toward HD\\,5980 (their Table~1) and conclude that a SNR shock is the most likely production mechanism for these ions in the observed quantities.  \\citet{Koenigsberger01} presented data on other ions not available to {\\it FUSE}.  In Figure~5, we present the observed column depths of these FUV ions and the values predicted by the models with v$\\rm_{s}$=160, 180 and 200 \\kms.  The ions are ordered according to the ionization potential needed to produce them.  The relative column depths of all the ions down to \\FeIII\\ are well matched by the 160 \\kms\\ model.  The higher velocity models predict too much \\OVI\\ relative to the other high ionization ions.  The column depths scale linearly with preshock density (all the models were run with n$_0$ = 1 cm$^{-3}$).  The offset between the 160 \\kms\\ model and the observations is $0.78\\pm0.14$ dex, so a preshock density of $\\rm n_o=6\\pm2 cm^{-3}$ can reproduce the observed column depths for the high ionization lines.  The shock velocity and preshock density toward HD\\,5980 are thus in good agreement with our measurements for Position~1.  The behavior of the low ionization lines is drastically different.  For \\SiII\\ and \\FeII, all the models predict column densities that are an order of magnitude higher than the observations (see Figure~5).  This dearth of low-ionization material can also be seen by comparing the observed and calculated ratio of \\OVI\\ and H$\\alpha$.  The shock models (which have been calculated to the point of complete recombination) predict I(\\OVI)/I(\\Ha)=3.3 for $\\rm v_s$=160 \\kms\\ and SMC abundances.  The observed ratio at Position~1 is  \\[  \\frac{I(OVI)}{I(H\\alpha)} = \\frac{SB(OVI)}{SB(H\\alpha)}~\\frac{C_{1035}}{C_{6563}}~\\frac{f_{H\\alpha}}{f_{OVI}} \\]  \\noindent where $SB$ indicates the surface brightness and $f$ the aperture filling factor of the relevant emission.  We assume the reddening correction $C_{1035}$=4.5 as discussed above; reddening at \\Ha\\ is moderate and we adopt $C_{6563}$=1.24.  The filling factor $f$ for each aperture is more uncertain.  We assume $f_{OVI}\\approx\\frac{1}{3}$ as discussed above.  It is difficult to say where emission does and does not appear at a fine scale in the low S/N echelle data.  The \\Ha\\ aperture was chosen carefully to encompass only the faint emission in Figure~3, but we estimate that $f_{H\\alpha}$ is still $\\sim$$\\frac{2}{3}$.  These values give I(\\OVI)/I(\\Ha)$\\approx$18.  Uncertainties in the filling factors, reddening, and surface brightnesses probably make this ratio uncertain by a factor of 2, but the ratio is still significantly higher than the predicted ratio of 3.3.  The weakness of \\Ha\\ emission compared with \\OVI\\ at Position~1 suggests the presence of an incomplete shock.  There are many cases where incomplete shocks are observed in individual filaments or small portions of SNRs \\citep{Fesen82,Blair91,Danforth01}; however, these regions are all at sub-parsec scales.  The LWRS aperture at Position~1 is sampling a region of \\SNR\\ roughly 8 pc across.  Furthermore, the observed \\Ha\\ surface brightness does not vary strongly between different portions of the SNR implying that the conditions at Position~1 are typical of the remnant as a whole.  In Figure~3, we see weak SNR \\Ha\\ emission or none at all.  This is a large object with moderate shock velocities implying that it is a middle-aged remnant like the Cygnus Loop or the Vela SNR; most of the optical emission from these remnants comes from complete, radiative shocks and these remnants are bright in \\Ha.  If some portions of \\SNR\\ were recombinationally complete, we would expect to see \\Ha\\ emission $\\sim$10 times brighter than is observed.  Thus, it appears that there is a global effect preventing the cooler parts of the recombination flow from forming.  We argue that the gas at the back end of the post-shock flow is being prevented from recombining by the ionizing flux from the bright stars in N66. There are eleven stars of spectral type O6.5 or earlier located 2\\arcmin\\ (35 pc, in projection) to the west of \\SNR.  The combined ionizing flux from NGC\\,346 is 4.0$\\times$10$^{40}$ erg s$^{-1}$ \\citep{Relano02} or $\\sim$1.5$\\times$10$^{51}$ ionizing photons s$^{-1}$.  HD\\,5980 itself contributes $\\sim4\\times10^{50}$ ionizing photons s$^{-1}$. Using basic ionization/recombination equations from \\citet{Osterbrock89} we find that this photon flux in a medium with density $\\sim$5 cm$^{-3}$ will yield an ionized sphere over 100 pc in radius.  Indeed, \\citet{Relano02} find that N66 is density bounded and that 45\\% of the ionizing photons from cluster stars escape to ionize the general SMC ISM.  This zone of ionization could encompass \\SNR\\ and affect the ionization balance in the recombining post-shock gas.  The dynamical time scale of the shock is of order a few hundred years.  The photoionization time scales are much shorter ($\\sim$100 days), while the recombination times are longer (of order 10,000 yr).  Thus the ionizing flux will suppress \\Ha\\ emission.  The strength and spectral energy distribution of the local ionizing flux will affect the column densities of the low and moderately ionized species.  In theory, much of the discrepancy between the modeled and observed \\FeII\\ and \\SiII\\ columns towards HD\\,5980 (Figure~5) could be due to iron and silicon being ionized to higher ionization stages.  If the best fit model is scaled up to n$\\rm_o=6$ cm$^{-3}$, the \\FeIII\\ column is underpredicted by about 0.3 dex.  This is clearly not enough to explain the 1.9 dex discrepancy in the \\FeII\\ column, but the stellar spectrum extends out to $\\sim$50 eV \\citep{Relano02} and the local flux could be producing \\ion{Fe}{4}.  However, no such discrepancies are seen in the \\SiIII\\ and \\SiIV\\ columns.  A more likely scenario is that the sight line to HD\\,5980 is passing through a small region where the swept up column is lower than average for \\SNR.  While the \\FeIII\\ zone is complete, the \\FeII\\ zone would be incomplete regardless of ionizing radiation.  A self-consistent shock model calculation that includes a photoionizing field is necessary to predict the column densities and the \\Ha\\ flux accurately.  While such modeling is beyond the scope of this paper, we conclude that the local ionizing field must have a significant effect on the observed characteristics of \\SNR\\ as a whole.  \\subsection{The Relationship Between \\SNR\\ and N66}  %  The emission and absorption data discussed here lead to a schematic structure for the N66--SNR region shown in Figure~6.  \\SNR\\ lies on the near side of N66 and is encountering the denser nebular material on its back side.  Position~1 is located at the edge of the X-ray emission (Figure~1).  The observed radial velocity of \\OVI\\ and \\Ha\\ emission at Position~1 is $\\sim$225 \\kms.  This is lower than the observed velocity of absorption ($+$300 \\kms) toward HD\\,5980 which lies projected closer to the center of the remnant.  Also it is $\\sim$75 \\kms\\ higher than the systemic velocity of N66 ($+$158 \\kms), which we would expect for material viewed tangentially at the edge of the expanding remnant.  Either the systemic velocity of the SNR and N66 differ by $\\sim$70 \\kms, or the observed X-ray edge (a portion of which is Position~1) does not represent the actual limb of the SNR shock.  The latter is assumed in Figure~6.  We propose that the X-ray emission arises over a spherical cap where the shock has encountered and heated denser material.  The remnant is not yet even half-way submerged in the N66 material, and thus the observed radius of X-ray emission is smaller than the actual blast wave radius.  The observed X-ray edge at Position~1 is inclined to our line of sight by $\\sim$60\\dd.  HD\\,5980 and the other UV-bright stars are located within N66 and have ionized a region of surrounding material.  The combined stellar winds and thermal pressure have formed a slowly expanding envelope around the cluster core.  Since high-velocity absorption is not seen in the spectrum of Sk80 \\citep{Hoopes01}, it must lie within or in front of the SNR.  HD\\,5980 lies behind the remnant and affects the ionization balance in the swept-up gas, preventing the formation of \\HI, \\SII, \\FeII, \\SiII\\ and other low-excitation species.  The bright nebular emission seen surrounding the UV bright stars of N66 arises in dense material behind an ionized layer.  The approaching SNR material is apparent in \\Ha\\ in Figure~3, but has not been conclusively identified in FUV spectra.  If the far side of the remnant is seen near $\\sim$300 \\kms\\ and if the SNR is centered at the SMC systemic velocity of $\\sim$155 \\kms, we expect the nearside absorption to be seen near v=10 \\kms.  This region is essentially masked by Galactic absorption making it difficult to detect the SNR.  \\citet{Hoopes01} compared the \\OVI\\ spectra of Sk80 and HD\\,5980 near zero velocity and found no difference.  However, both of these sight lines are toward the remnant and may contain nearly identical foreground \\OVI\\ columns.  The four {\\it FUSE} targets in NGC\\,346 lie less than 2\\arcmin\\ from HD\\,5980 and are outside the SNR sight line.  A comparison of the averaged, normalized \\OVI\\ profiles from these four targets and the two sight lines projected within the SNR is shown in Figure~7.  We find a small but systematic excess in \\OVI\\ absorption near v=0 \\kms\\ in the two remnant profiles.  The difference represents an extra \\OVI\\ column density of about $10^{13.8}$ cm$^{-2}$ or about $\\frac{1}{4}$ that of the backside value reported by \\citet{Hoopes01}.  We also find that the velocity centroid of the HD\\,5980/Sk80 zero velocity absorption is shifted by about $+$10 \\kms\\ with respect to the NGC\\,346 Galactic absorption.  This is not too different from the velocity found by \\citet{Koenigsberger01} for \\SiIV\\ and \\CIV\\ absorption features that they attribute to front-side SNR material.  Recent work by \\citet{Howk02} on Galactic halo \\OVI\\ absorption has shown significant variations in \\OVI\\ column over as little as 3\\arcmin.  The variation in N(\\OVI) between the four NGC\\,346 targets (mean separation 19\\arcsec) is 9\\% (the maximum variation between two measurements is 12\\%) while the variation between the two SNR sight lines (separation $<$1\\arcmin) is only 8\\%, comparable to the uncertainties due to continuum placement, noise, and limits of integration.  However, the cluster and SNR stars show a relative difference of 53\\% in the zero-velocity \\OVI\\ column over $\\sim$2\\arcmin.  Because of the potential Galactic halo variation on small scales, the matter of frontside absorption remains inconclusive.  However, it is plausible we are seeing front side absorption near v=$+$10 \\kms\\ in the HD\\,5980/Sk80 data.  If true, then Sk80 would be inside \\SNR, as shown in Figure~6.  Even if it is not, we can place an upper limit of N(\\OVI)$\\la6\\times10^{13}$ cm$^{-2}$ for the frontside \\OVI\\ column density, much less than is observed from the back side of \\SNR.  In a recent paper, \\citet{Velazquez03} model the collision of a SNR shock with a wind-blown bubble around an evolved star and suggest that the far side of \\SNR\\ is interacting with the stellar wind of HD\\,5980.  The X-ray image of the system \\citep{Xmega} does, in fact, show some resemblance to the models and there is reasonable agreement between predicted and observed X-ray fluxes.  Our observations allow us to further test this wind-SNR model.  The Vel\\'azquez et~al. wind-blown bubble has a radius of $\\sim$10 pc while our Position~1 (at the rim of the SNR) is at least twice that distance (in projection) from HD\\,5980; thus, Position~1 is sampling the ISM, not stellar wind.  In \\S4.1 we derived a preshock density of $\\sim$6 cm$^{-3}$ for the material at both Position~1 and in the shock toward HD\\,5980, at least two orders of magnitude higher than the interior of the modeled wind bubbles \\citep[see Figures~1 and 2 of][]{Velazquez03}.  Furthermore, the $+$300 \\kms\\ component exhibits an ionization structure expected for a single SNR shock \\citep{Hoopes01}.  If this component arose in a turbulent region between SNR and stellar wind shock fronts, as suggested by \\citet{Velazquez03}, we might not expect to see the kinematics structure observed by \\citet{Hoopes01}.  Clearly, higher resolution spectral observations and detailed shock models are required.  \\subsection{Implications for SNRs in \\HII\\ Regions}  %  Core-collapse supernovae arise from massive stars often formed in clusters and associations--the same massive stars which photoionize their surrounding gas and create bright \\HII\\ regions.  Massive SN progenitors cannot migrate far from their birth places in their short lifetimes and thus SNRs associated with H~II regions should occur frequently.  To various degrees, this is not seen in optical SNR surveys of nearby galaxies.  The primary optical method for detecting SNRs, especially in galaxies beyond the Magellanic Clouds, is to compare \\Ha\\ to [\\SII] emission, and pick out objects with [\\SII]/\\Ha $\\ge$ 0.4 \\citep[cf.][and references therein]{Long90,BL97}.  In recent years, as the sensitivity and resolution of radio and X-ray observations has become competitive with optical ground-based observations, independent searches at these wavelengths have been undertaken and compared against the optical SNR lists \\citep{Lacey97,Gordon98,Gordon99,Pan00,Pan02}.  These multiwavelength surveys and comparisons provide a more complete survey method, but in general do not find many of the same SNRs.  Discussions of selection effects usually point toward ``confusion'' as being responsible for incompleteness of optical SNR surveys in crowded and complex regions of emission.  Clearly this is plausible at some level, but the denser conditions in \\HII\\ regions might also be expected to cause brighter optical (radiative) SNRs since the shock emissivity increases as the density squared.  Scanning through the optical SNR catalogs of nearby galaxies \\citep[to name a few]{BL97,Long90,MF97}, one can find a number of stunning counter examples to the ``confusion'' hypothesis.  There are often bright optical SNRs found in confused regions and adjacent to bright \\HII\\ region emission.  The mechanism we propose in \\S4.2 seems to be a more plausible explanation for why the optical SNR searches miss some SNRs and not others in confused regions.  If a supernova occurs close enough to ionizing stars, the potential bright optical SNR emission in \\Ha\\ and \\SII\\ is prevented from forming and no optical SNR can be detected, especially against bright nebular emission.  Our investigation of \\SNR, however, demonstrates that bright FUV emission lines are still present in such a situation and may provide an alternate way of not only detecting but studying such SNRs in detail.  Detecting SNRs with higher ions circumvents the problems of truncated recombination and bright background emission.  Stellar photoionization will not prevent the production of high ions nor will they be contaminated significantly by surrounding nebular emission.  Our observations have shown that in a bright \\HII\\ region, the detection of localized emission from high-ionization species such as \\OVI\\ can provide a clear signature of SNR emission.  Position~1 lies on the edge of the SNR and shows strong emission in \\OVI\\ and \\CIII.  Position~2 lies outside the remnant but still within bright optical emission; yet no line emission is present.  In conjunction with radio and X-ray data, such a detection is conclusive.  We note however that the use of the FUV is restricted to lines of sight where extinction is low.  Nearby extragalactic targets such as the Magellanic Clouds and Local Group galaxies will be the easiest to investigate with this technique."
    },
    "0212/astro-ph0212024_arXiv.txt": {
        "abstract": "Experience with ASTROVIRTEL, scientific analysis of current large data sets, and detailed preparation for the truly huge future missions especially GAIA, provide important lessons for the Astrophysical Virtual Observatory. They demonstrate that the science cases are impressive, specifically allowing new thresholds to be crossed. The AVO is more than just faster cheaper better, it allows the new.  The example of use of pre-explosion imaging of supernova to identify progenitors is used to illustrate some general challenges.  Some non-trivial technical astronomical issues arise, especially astrometry, to complement the many technical implementation challenges.  A critical scientific lesson is the need to quantify data quality. How are we to ensure the Virtual Observatory produces top science, and avoids being overwhelmed with mediocre data? ",
        "introduction": "The Astrophysical Virtual Observatory has an exceptionally strong science case. One example, identification of the pre-explosion state of core-collapse supernovae, is described below. This case illustrates some of the technical and strategic challenges which the Virtual Observatory projects have still to face.  There is a strong scientific case to identify and access appropriate archival data on the sites of supernova explosions. Since the candidate star is identified by its self-immolation, the only relevant data are archival!  A project to obtain suitable new data to act as a future archive is underway with HST. In addition, much valuable data already exists. Some of this is in archives, and is well calibrated. Much is in private hands. The effort needed to use even excellent quality calibrated and published data is illustrated below, using as examples the searches by Smartt etal for the progenitors of SN1999em and SN2002ap.   One the most difficult of the astronomical challenges facing the integrated use of federated multi-wavelength multi-resolution archives involves source extraction, and astrometric cross-matching. Some examples are given below. The challenge is however obvious to all who have experience even with combination of HST WFPC2 and NICMOS images. Even in this case, with high-quality, high spatial resolution, stable, well-calibrated data sets, with only a one-half decade wavelength range, simple matching of optical and near-IR images of star clusters is not trivial. A nice example is available in figures one and two of Johnson etal (2001), who present PC and NICMOS images of two young LMC globular clusters.  A second challenge involves data reliability. The example below illustrates how even unusually high-quality data cannot be used beyond the range in which their systematic uncertainties become relevant to the science at hand. However, few (if any) data archives are calibrated well enough to provide this information, except in response to very specific questions and applications.  This raises the spectre of well-meaning providers of what are in fact data of limited calibration ensuring that virtual observatory data product users either produce defective science, or are overwhelmed with learning the limits of every individual data set accessed by the entire system. Might it be that data-archives need a quality-assessment check before they are `eligible' for access? Or is it to be {\\em caveat emptor?} ",
        "conclusions": "Identification and analysis of high-resolution ground-based images of the pre-explosion sites of Type\\,II supernovae provide unique information on the late stages of evolution of massive stars, the chemical evolution of the Universe, and the physics of feedback on galaxy evolution. Examples of the successful application of this method are presented above. From a Virtual Observatory perspective there are some important lessons to be learned:  \\begin{itemize} \\item Even excellent quality and carefully derived data products (crowded field photometry in the specific example here) can prove unsuitable for purposes different than their original application. The scientific integrity of data retrieved by the Virtual Observatory `system' must still be established by the responsible astronomer. \\item A considerable amount of non-archived non-calibrated data exist in private repositories. While in some cases this can be of unique value, the work involved in retrieving, and especially in calibrating, old data suggests this is a worthwhile use of resources only in special cases. \\item It is interesting to consider if the point above means that the AVO should restrict itself to accessing only well-described and major public data sets. \\item A general challenge is to provide adequate astrometric cross-matching for different datasets. This raises the issue of different spatial resolutions, source dropouts, etc. The examples given here suggest that no complete general solution is feasible, except for imaging data from telescopes with extremely well-quantified optical systems. \\end{itemize}"
    },
    "0212/astro-ph0212212_arXiv.txt": {
        "abstract": "{We report the discovery, based on an inspection of the OGLE-II database, of a group of blue variables in the Magellanic Clouds showing simultaneously two kinds of photometric variability: a short-term cyclic variability with typical amplitude $\\Delta I \\sim$ 0.05 mag and period $P_{1}$ between 4 and 16 days and a sinusoidal, long-term cyclic oscillation with much larger amplitude $\\Delta I \\sim$ 0.2 mag with period  $P_{2}$ in the range of 150-1000 days. We find that both periods seems to be coupled through the relationship $P_{2}$ = 35.2 $\\pm$ 0.8 $P_{1}$. In general, the short term variability is reminiscent of those shown by Algol-type binaries. We propose that the long-term oscillation could arise in the precession of a elliptical disc fed by a Roche-lobe filling companion in a low mass ratio Algol system. ",
        "introduction": "Over the past years, the microlensing projects (OGLE, MACHO, EROS) have monitored millions of stars in the Magellanic Clouds and Galactic bulge for variability. The resulting huge photometric databases are very well suited not only for microlensing studies but also for many other issues of modern astrophysics, including the distance scale, variable stars, star clusters etc. In particular, the OGLE-II project (Udalski, Kubiak and Szymanski, 1997), has provided accurate BVI measurements for about 6.5 million stars from the central parts of the Magellanic Clouds (Udalski et al.\\ 1998, 2000). Based on this same material, a unique catalog containing about 68.000 variable stars has just been released ({\\.Z}ebru{\\'n} et al.\\ 2001).  One example of the scientific information that it is possible to extract from such a catalogue is shown by the work of Mennickent et al. (2002, hereafter M02), who report the discovery of Be star candidates in the Small Magellanic Cloud showing unexpected photometric variations. Basically, these authors found four types of variability in targets within the luminosity-colour box of typical Be stars: Type-1 stars showing outbursts, Type-2 stars showing sudden magnitude jumps, Type-3 stars showing periodic variations and Type-4 stars showing random variability. According to these authors, possible causes for Type-1 and Type-2 stars include blue pre-main sequence stars surrounded by thermally unstable accretion discs and white dwarfs accreting from the Be star envelope in a Be+WD binary. Type-4 stars could be classical Be stars. On the other hand, Type-3 stars could not be linked to the Be star phenomenon at all, mainly due to their rather red colours and strictly periodic behaviour. The above study is complemented by the work by Keller et al. (2002) on blue variable stars from the MACHO database in the LMC. These authors report basically the same photometric behaviour in a sample of LMC blue variables. These authors also find emission lines in the spectra of many of these variables.  Further inspection of the misterious Type-3 stars found in the SMC and the LMC (Mennickent et al. in preparation) revealed 30 objects showing double-periodic photometric variations: a long-period sinusoidal variation and a short-period modulation. Some of these  short-term light curves are typical of eclipsing Algol variables. The study of this new double-period photometric variability is the subject of this paper. ",
        "conclusions": "The long-term well behaved sinusoidal modulations observed in our sample is not known to occur in other types of blue variable stars. Here we propose as a possible cause for the long-term oscillation the precession of an elliptical disc around a blue star in a semi-detached binary system. The origin of the disc ellipticity and precession could be the tidal interaction between the disc and the Roche-lobe filling secondary star, in a similar way that occurs in short-orbital period dwarf novae of the SU UMa class during superoutburst (e.g. Patterson 2001).   \\begin{table*} \\caption[]{The double-periodic blue stars. OGLE name and MACHO identification are given for every star, along with their photometric period and error. The half amplitude of the long-term variation is also given for the OGLE $I$ band and the MACHO $b$ band, along with the epoch (HJD$-$2\\,400\\,000) for the faint-to-bright mean level crossing for the MACHO $b$ band. The half amplitude of the MACHO $b-r$ light curve is also given, along with their difference in phase with the $b$ curve, considering the epoch for positive to negative mean $b-r$ level crossing as reference. Dashes indicate non-detected or marginal colour variations. A note indicates the appearance of the short-term variability: double-wave (dw), single wave (sw) or eclipsing with two minima (e).} \\begin{center} \\begin{tabular}{lrrrrrrrrr} \\hline  \\multicolumn{1}{c}{Star}& \\multicolumn{1}{c}{MACHO ID} & \\multicolumn{1}{c}{$P_{1}$ (d)} & \\multicolumn{1}{c}{$P_{2}$ (d)} & \\multicolumn{1}{c}{$A_{I}$}   & \\multicolumn{1}{c}{$A_{b}$}   & \\multicolumn{1}{c}{$E_{b}$}   & \\multicolumn{1}{c}{$A_{b-r}$}   & \\multicolumn{1}{c}{$\\Delta\\Phi$}   & \\multicolumn{1}{c}{Note}\\\\ \\hline \\hline OGLE00451755-7323436 &212.15675.158  &  5.178(5)   &  171(15)  &0.12 &0.11&50350.2(8)&0.02&-0.35& dw   \\\\ OGLE00474820-7319061 &212.15847.466  &  5.497(7)   &  177(10)  &0.12 &0.10&50329.3(7)&- &-& dw      \\\\ OGLE00553643-7313019 & 211.16304.169 &  5.092(5)   &  176(17)  &0.08  &0.04&50238(1)&0.02&0.53&   dw        \\\\ OGLE05025323-6909493 &  1.3686.53    &  8.025(18)  &  255(30)  &0.17  &-&-&-&-&                 dw \\\\ OGLE05040378-6917508 &  1.3805.130   &  6.223(7)   &  207(20)  &0.11  &-&-&-&-&                 dw  \\\\ OGLE05060009-6855025 &  1.4174.42    &  3.849(7)   &  230(22)  &0.19  &-&-&-&-&                 sw     \\\\ OGLE05101621-6854290 &  79.4900.185  &  4.301(6)   &  319(34)  &0.06  &0.04&50344(2) &0.02&0.53  & sw     \\\\ OGLE05115466-6846369 &  2.5144.4555  &  9.138(2)   &  361(29)  &0.26  &-&-&-&-&                 e\\\\ OGLE05142677-6910559 &  79.5501.400  &  6.515(2)   &  224(14)  &0.13  &0.08&50430(2)&0.04&-0.43  & e  \\\\ OGLE05143758-6852259 &  79.5506.139  &  5.372(6)   &  185(11)  &0.14  &0.10&50237.3(5)&0.03&0.55 &dw?   \\\\ OGLE05152654-6923257 &  79.5740.5092 &  6.292(8)   &  205(15)  &0.04  &0.04&50275.3(6)&-&-& dw    \\\\ OGLE05155332-6925581 &  79.5739.5807 &  7.2835(16) &  188(11)  &0.11  &0.08&50254(1)&0.03&0.54 &  e      \\\\ OGLE05171401-6936374 &  78.5979.58   &  8.309(4)   &  311(21)  &0.13  &0.08&50345(1)&0.03&0.48 &  e  \\\\ OGLE05194110-6931171 &  78.6343.81   &  6.9044(10) &  226(13)  &0.11  &0.06&50247(2)&0.02&0.59  & e  \\\\ OGLE05195898-6917013 &  80.6468.83   &  2.410(10)  &  140(5)   &0.06  &0.04&50267.6(8)&-&-  &sw       \\\\ OGLE05203325-6910146 &  80.6469.95   &  5.737(5)   &  182(10)  &0.04  &0.05&50357.1(8)&-&- &dw        \\\\ OGLE05260516-6954534 &  77.7426.140  &  3.632(3)   &  233(15)  &0.07  &0.07&50425.7(7)&-&-  &     sw       \\\\ OGLE05274332-6950556 &  77.7669.1013 &  7.320(11)  &  227(20)  &0.05  &0.03&50280(1)&0.02&0.54 &  dw   \\\\ OGLE05285370-6952194 &  77.7911.26   &  15.854(31) &  620(70)  &0.07  &0.06&50388(2)&0.01&0.67  & dw    \\\\ OGLE05294913-6949103 &  77.8033.140  &  7.184(8)   &  258(20)  &0.10  &0.09&50382.2(6)&0.01&-0.38& dw   \\\\ OGLE05295881-6934075 &  77.8036.5142 &  5.597(4)   &  179(8)   &0.12  &-&-&-&-&                 dw     \\\\ OGLE05313130-7012584 &  7.8269.36    &  9.231(21)  &  960(176) &0.04  &0.03&51226(4)&-&-  & sw         \\\\ OGLE05333926-6956229 &  81.8636.51   &  7.863(15)  &  257(18)  &0.03  &0.01&50335(2)&0.01&0.52 &  dw         \\\\ OGLE05371342-7010580 &  11.9237.2121 &  10.913(23) &  421(40)  &0.12  &-&-&-&-&                 dw         \\\\ OGLE05390681-7027487 &  11.9475.96   &  6.967(12)  &  276(15)  &0.10  &0.01&50335(2)&0.01&0.52   &dw      \\\\ OGLE05390992-7019262 &  11.9477.138  &  6.632(8)   &  198(15)  &0.10  &0.05&50301(1)&0.02&0.56  & dw      \\\\ OGLE05391746-7044019 &  11.9592.22   &  7.151(8)   &  219(20)  &0.06  &0.04&50401.0(8)&0.01&-0.49& dw        \\\\ OGLE05410217-7011043 &  76.9842.2444 &  7.352(13)  &  264(26)  &0.11  &0.11&50395(2)&-&-   &      dw       \\\\ OGLE05410942-7002215 &  76.9844.110  &  6.586(9)   &  245(23)  &0.12  &0.09&50510(2)&-&-     &    dw         \\\\ OGLE05435003-7057431 &  15.10314.144 &  5.012(5)   &  173(12)  &0.12  &0.09&50250.7(7)&-&-    &   dw         \\\\ \\hline \\end{tabular} \\end{center} \\end{table*}     The 3:1 resonance between a disc particle orbiting the primary and the binary system occurs at $R_{disc} \\approx 0.46a$, where $a$ is the binary separation. Elliptical orbits at this radius will experience a dynamical apsidal advance with a period given by:\\\\  $P_{p} = \\frac{\\sqrt{1+q}}{0.37q}(R_{disc}/0.46a)^{-2.3}P_{o}$ \\hfill(3)\\\\  (Murray 2000) where $q$ is the ratio between the mass of the secondary star and the primary star and $P_{o}$ the orbital period. If we interpret $P_{1}$ as the orbital period  and $P_{2}$ as the precession period, then the above equation imposes a strong correlation between both periodicities, as effectively seen in Fig.\\,3. The observed slope of $P_{2}/P_{1}$ $\\approx$ 35 should imply a mass ratio $q \\approx$ 0.08 for $R_{disc} = 0.46 a$. There are few Algols with $q <$ 0.1 (e.g. Richards \\& Albright 1999), and they do not show the long-term oscillations described in this paper. However it is exciting that only in these low $q$ systems two important conditions are fulfilled: the primary radius is smaller than the stream circularization radius, being possible the formation of an accretion disc (e.g. Richards \\& Albright 1999) {\\it and} the 3:1 resonance radius is below the tidal truncation radius, so the disc could in principle grown beyond the 3:1 resonance and may eventually start to precess (Fig.\\,4). Long-term spectroscopic observations and detailed modeling are required to confirm this scenario."
    },
    "0212/astro-ph0212097_arXiv.txt": {
        "abstract": "{ We present cosmological parameter constraints on flat cosmologies dominated by dark energy using various cosmological data including the recent Archeops angular power spectrum measurements. A likelihood analysis of the existing Cosmic Microwave Background data shows that the presence of dark energy is not requested, in the absence of further prior.  This comes from the fact that there exist degeneracies among the various cosmological parameters constrained by the Cosmic Microwave Background.  We found that there is a degeneracy in a combination of the Hubble parameter $H_0$ and of the dark energy equation of state parameter $w_\\QUINT$, but that $w_\\QUINT$ is not correlated with the primordial index $n$ of scalar fluctuations and the baryon content $\\Omega_\\BAR h^2$. Preferred primordial index is $n = 0.95 \\pm 0.05 (68\\%)$ and baryon content $\\Omega_\\BAR h^2 = 0.021 \\pm 0.003$.  Adding constraint on the amplitude of matter fluctuations on small scales, $\\sig$, obtained from clusters abundance or weak lensing data may allow to break the degenaracies, although present-day systematics uncertainties do not allow firm conclusions yet. The further addition of the Hubble Space Telescope measurements of the local distance scale and of the high redshift supernovae data allows to obtain tight constraints. When these constraints are combined together we find that the amount of dark energy is $0.7^{+0.10}_{-0.07}$ ($95\\%$ C.L.) and that its equation of state is very close to those of the vacuum: $w_\\QUINT < -0.75 $ ($> 95\\%$ C.L.).  In no case do we find that quintessence is prefered over the classical cosmological constant, although robust data on $\\sig$ might rapidly bring light on this important issue. ",
        "introduction": "The determination of cosmological parameters has always been a central question in cosmology. In this respect the measurements of the Cosmological Microwave Background (CMB) anisotropies on degree angular scales has brought one of the most spectacular results in the field: the flatness of the spatial geometry of the universe, implying that its density is close to the critical density.  Although, during the last twenty years the evidence for the existence of non-baryonic dark matter has strongly gained in robustness, observations clearly favor a relatively low matter content somewhere between $20$ and $50\\%$ of the critical density, indicating that the dominant form of the density of the universe is an unclustered form. Furthermore, the observations of distant supernovae, at cosmological distance, provide a direct evidence for an accelerating universe, which is naturally explained by the gravitational domination of a component with a relatively large negative pressure, $P_\\QUINT = w_\\QUINT \\rho_\\QUINT$ with $w_\\QUINT < - 1 / 3$.  The cosmological constant $\\Lambda$ (for which $w_\\Lambda = - 1$) is historically the first possibility which has been introduced and which satisfies this requirement.  However, the presence of a non-zero cosmological constant is a huge problem in physics: (i) quantum field theory predicts that $\\Lambda$ should be the sum of a number of enormous contributions, so in order to avoid a cosmological catastrophe, it is usually assumed that a yet unknown mechanism produces a cancelation between all these contributions; (ii) it is difficult to think of a mechanism which puts $\\Lambda$ to $0$ exactly, but it is even more difficult to find a mechanism which gives $\\rho_\\Lambda \\sim \\rho_\\CRIT \\sim 10^{- 122} \\rho_\\PL$, as the supernovae observations suggest, where $\\rho_\\PL$ and $\\rho_\\CRIT$ are the Planck energy density and the critical density today, respectively. For this reason the concept of quintessence, a scalar field with negative pressure, has recently been proposed as a possible alternative to a cosmological constant.  In this paper we shortly describe the quintessence paradigm and its effect on some observable quantities. We then summarise the different sets of data and methods used to constrain cosmological parameters. We then conclude by showing the results on quintessence and cosmological parameters. ",
        "conclusions": "We have studied the constraints that can be obtained on cosmological parameters within the quintessence paradigm by using various combinations of observational data set. For simplicity, only models with constant $w_\\QUINT$ were examined as they are believed to be sufficient to existing data at their present accuracy level. For similar reasons, we neglected possible reionization or non-zero gravitational wave contribution.  Indeed, this approach has the advantage to help to understand whether quintessence models should be favored or not over a classical cosmological constant, while constraints on more specific scenarios have to be investigated directly one by one (Douspis \\ETAL{} \\cite{dousp2}). Our analysis method has been to investigate contours in 2-D parameters space. Such an approach allows to examine possible degeneracies among parameters which are not easy to identify when constraints are formulated in term of a single parameter. For instance, we found that CMB data alone, despite the high precision data obtained by Archeops do not require the existence of a non-zero contribution of quintessence, because of the degeneracy with the Hubble parameter: in practice CMB data leave a large fraction of the $(\\Omega_\\QUINT, w_\\QUINT)$ plane unconstrained, while only a restricted region of the $(\\Omega_\\QUINT, H_0)$ plane is allowed. On the contrary, we found that almost no correlation exist with the baryonic content $\\Omega_\\BAR$ nor the primordial index $n$. In order to restrict the parameter space of allowed models, we have applied several different constraints. Interestingly, we found that the amplitude of the dark matter fluctuations, as measured by clusters abundance or large scale weak lensing data can potentially help to break existing degeneracies, although existing uncertainties, mainly systematics in nature, do not allow firm conclusion yet. Clearly this will be an important check of consistency in the future. We have then added constraints from supernovae data as well as HST estimation of the Hubble parameter in order to break existing degeneracies. This allows us to infer very tight constraints on the possible range of equation of state of the dark energy. Probably the most remarkable result is that no preference for quintessence does emerge from existing CMB data although, accurate measurement of the amplitude of matter fluctuations on scale of $8 \\, h^{-1} \\UUNIT{Mpc}{}$ may change this picture."
    },
    "0212/astro-ph0212268_arXiv.txt": {
        "abstract": "Collapse of the rotating spheroid is approximated by a system of ordinary differential equations describing its dynamics. The gravitational potential is approximated by the one of the iniform Maclaurin spheroid. Developement of gravitational instability and collapse in the dark matter medium do not lead to any shock formation or radiation, but is characterized by non-collisional relaxation, which is accompanied by the mass and angular momentum losses. Phenomenological account of these processes is done in this model. Formation of the equilibrium configuration dynamics of collapse is investigated for several parameters, characterizing the configuration. A very long gravitational wave emission during the collapse is estimated, and their possible connection with the observed gravitational lenses is discussed. ",
        "introduction": "A study of of a formation of dark matter objects in the universe is based on N-body simulations, which are very time consuming. In this situation a simplified approach may become useful, because it permits to investigate rapidly by analytical or by simple numerical calculations many different variants , and obtain crudly some new principally important features of the problem which could be lost or not visible during a long numerical  work.  The modern theory of a large scale structure is based on the ideas of Zeldovich (1970) about a formation of strongly non-spherical structures during non-linear stages of a development of the gravitational instability, known as \"Zeldovich's pancakes\". The numerical simulations started by Doroshkevich et al. (1980) had been performed subsequently by many groups, revealing complicated structures of the new born objects (see e.g. Doroshkevich et al., 1999).  Here we derive and solve equations for a simple model of a dynamical behaviour of a compressible rotating spheroid in which a motion along both axies takes place in the common gravitational field of a uniform spheroid, and under the action of the isotropic pressure (taken into account in the approximate non-differential way), and otherwise independently on each other. Due to non-interacting particles of the dark matter a strong compression does not lead to a formation of the shock wave and additional energy losses due to radiation at increasing of a temperature.. All losses are connected with a runaway particles, and relaxation consists only in the phase mixing leading effectively to a transformation of the kinetic energy of the ordered motion (collapse) into the kinetic energy of the chaotic motion, and to creation an effective pressure and thermal energy. In presence of the chaotic motion the pancake formed during the collapse has a finite minimal thickness, because particles cross the equatorial plane non-simultaneously, forming therefore the effective pressure, which \"stops\" the contraction. These effects may be described approximately in the hydrodynamical approach, in which there are no shock waves and the main transport process is an effective bulk viscosity, leading to relaxation and to a damping of the ordered motion. This relaxation is based on the idea of a \"violent relaxation\" of Lynden-Bell (1967). The system of the ordinary differential equations is derived, which describes the dynamical behaviour of the compressible spheroid, where the relaxation and losses of the energy, mass, and angular momentum are taken into accound phenomenologically. In absence of any dissipation these equations describe a self-consistent conservative system, where gravitational and thermal pressure forces are present. Namely, the equilibrium solution of these equation describes a Maclaurin spheriod, where the density is not prescribed but is found self-consistently, when the effective entropy is known. We consider here only non-relativistic dark matter with a relations between the density $\\rho$, energy density $E$, pressure $P$, and specific entropy $S$ as  $$\\rho E =\\frac{3}{2}P, \\quad P\\sim \\rho^{5/3}, \\quad E\\sim \\rho^{2/3} \\quad {\\rm at\\,\\,\\, constant}\\,\\,\\, S. $$ Solution of simple ordinary equations reveals some interesting features of the collapse of non-interacting matter, which had not been noticed earlier. The collapse of a spherioid, leading to the formation of a pancake may be continued by the fornation of a transient oblate figure, strongly elongated along the axis of a symmetry with a ratio of the axis $c$ to the equatorial radius $a$ : $c/a \\sim 5$. Formation of such transient figures, appearing only at low rotation, could be expected in central parts of the collapsing dark matter objects during initial stages of the collapse, when the influence of the outer non-uniform part of the cloud is still negligible.  We have found, that at a realistic rate of the relaxation rate with a characteristic time equal to 3 local Jeans times, the oscillations do not damp completely, and about 1/10 of the initial amplitude of a velocity survives after at least 10 oscillationl periods. That means that the most massive dark matter objects may be still in the oscillatory state. These oscillations, as well as variable gravitational fields of the collapsing dark matter objects could be important for the the interpretation of the observed picture of the cosmic microvave background (CMB) fluctuations. Very long gravitational waves emitted mainly during the first stage of the pancake formation (see Thuan and Ostriker, 1974; Novikov, 1976) could be also important for the CMB fluctuations (Doroshkevich et al., 1977), as well as for the gravitational lensing of the distant objects. ",
        "conclusions": "The simplified model of the collapse and subsequent relaxation of the rotating compressed shperoid may be related to the processes in the central parts ob the dark matter objects. The formation of the transient prolate dark matter spheroid follows from the calculations for slowly rotating (or non-rotating) objects.  For realistic values of the relaxation time we have obtained rather slow damping of the oscillations, indicating that after $\\sim 10$ oscillations the amplitude may remain on the level of 1/10 of the Hubble value of the velosity for a corresponding mass scale. Such oscillations in the dark matter objects of the largest scale may be preserved until the present time.  Interaction of the cosmic microvawe background radiation (CMB) with the dark matter objects on the stage of their collapse and subsequent oscillation phases may influence the visible characteristics of the observed picture of the fluctuations of CMB.  The weak but very long gravitationl waves (GW) emitted mainly on the stages of the collapse and pancake formation form a long wave GW background, which also influence CMB fluctuations (Doroshkevich et al, 1977).  Besides, the existence of such a long GW in the space between the source and the observer may be registered as an action of the gravitational lense. Absence of the appropriate lensing objects in the direction of any observed gravitational lense object may indicate to the existence of dark matter objects without the barion matter presence, or may be the indication to the presence of the very long GW which is moving in the space between us and the lensed object The expected angles of such lensing ($\\sim 10^{-3}$) arc sec. may be available in the near future.  The present approach describing the compressed spheroid by the two degrees of freedom may be evidently generalized for a description of the 3-axis spheroid by the ordinary equations with 3 degrees of freedom. In particular, it would be possible to investigate the influence of the damping processes on the boundaries of the stability of the Maclaurin spheroid for the transformation into the Jacobi 3-axis ellipsoid (secular and dynamical) at large angular momentums. The model could be generalized also for the case of anizotropic relaxation and creation of anizotropic pressure.  \\begin{figure} \\centerline { \\psfig{figure=figu1.eps,width=11cm,angle=-90} } \\caption{Time dependence of semiaxies $a(t)$ and $c(t)$, and sinus of the rotational phase: $phase(t)$ for initial values of the variant 1, Table 1. } \\label{fig1} \\end{figure}  \\begin{figure} \\centerline{ \\psfig{file=figu1a.eps,width=10cm,angle=-90} } \\caption{Time dependence of the non-dimensional entropy ${\\cal E}(t)$, mass $m(t)$, and total energy $U_{tot}(t)$ for initial values of the variant 1, Table 1.} \\label{fig1a} \\end{figure}  \\begin{figure} \\centerline { \\psfig{figure=figu2.eps,width=11cm,angle=-90} } \\caption{Same as in Fig.1, for initial values of the variant 2, Table 1. } \\label{fig2} \\end{figure}  \\begin{figure} \\centerline{ \\psfig{file=figu2a.eps,width=10cm,angle=-90} } \\caption{Same as in Fig.2, for initial values of the variant 2, Table 1.} \\label{fig2a} \\end{figure}  \\begin{figure} \\centerline { \\psfig{figure=figu3.eps,width=11cm,angle=-90} } \\caption{Same as in Fig.1, for initial values of the variant 3, Table 1. } \\label{fig3} \\end{figure}  \\begin{figure} \\centerline{ \\psfig{file=figu3a.eps,width=10cm,angle=-90} } \\caption{Same as in Fig.2, for initial values of the variant 3, Table 1.} \\label{fig3a} \\end{figure}  \\begin{figure} \\centerline { \\psfig{figure=figu4.eps,width=11cm,angle=-90} } \\caption{Same as in Fig.1, for initial values of the variant 4, Table 1. } \\label{fig4} \\end{figure}  \\begin{figure} \\centerline{ \\psfig{file=figu4a.eps,width=10cm,angle=-90} } \\caption{Same as in Fig.2, for initial values of the variant 4, Table 1.} \\label{fig4a} \\end{figure}  \\begin{figure} \\centerline { \\psfig{figure=figu5.eps,width=11cm,angle=-90} } \\caption{Same as in Fig.1, for initial values of the variant 5, Table 1. } \\label{fig5} \\end{figure}  \\begin{figure} \\centerline{ \\psfig{file=figu5a.eps,width=10cm,angle=-90} } \\caption{Same as in Fig.2, for initial values of the variant 5, Table 1.} \\label{fig5a} \\end{figure}  \\begin{figure} \\centerline { \\psfig{figure=figu6.eps,width=11cm,angle=-90} } \\caption{Same as in Fig.1, for initial values of the variant 6, Table 1. } \\label{fig6} \\end{figure}  \\begin{figure} \\centerline{ \\psfig{file=figu6a.eps,width=10cm,angle=-90} } \\caption{Same as in Fig.2, for initial values of the variant 6, Table 1.} \\label{fig6a} \\end{figure}  \\begin{figure} \\centerline { \\psfig{figure=figu7.eps,width=11cm,angle=-90} } \\caption{Same as in Fig.1, for initial values of the variant 7, Table 1. } \\label{fig7} \\end{figure}  \\begin{figure} \\centerline{ \\psfig{file=figu7a.eps,width=10cm,angle=-90} } \\caption{Same as in Fig.2, for initial values of the variant 7, Table 1.} \\label{fig7a} \\end{figure}  \\begin{figure} \\centerline { \\psfig{figure=figu8.eps,width=11cm,angle=-90} } \\caption{Same as in Fig.1, for initial values of the variant 8, Table 1. } \\label{fig8} \\end{figure}  \\begin{figure} \\centerline{ \\psfig{file=figu8a.eps,width=10cm,angle=-90} } \\caption{Same as in Fig.2, for initial values of the variant 8, Table 1.} \\label{figa} \\end{figure}  \\medskip  {\\Large{\\bf Acknowledgement}}  \\medskip  I grateful to Theoretical Astrophysics Center (TAC), Copenhagen, for hospitality during the work on this paper, and to A.G. Doroshkevich and P.D. Naselsky for useful discussions."
    },
    "0212/astro-ph0212118_arXiv.txt": {
        "abstract": "\\chandra\\ X-ray observations of \\dem71, a supernova remnant (SNR) in the Large Magellanic Cloud (LMC), reveal a clear double shock morphology consisting of an outer blast wave shock surrounding a central bright region of reverse-shock heated ejecta.  The abundances of the outer shock are consistent with LMC values, while the ejecta region shows enhanced abundances of Si, Fe, and other species. However, oxygen is not enhanced in the ejecta; the Fe/O abundance ratio there is $\\gsim$5 times the solar ratio.  Based on the relative positions of the blast wave shock and the contact discontinuity in the context of SNR evolutionary models, we determine a total ejecta mass of $\\sim$1.5 $M_\\odot$. Ejecta mass estimates based on emission measures derived from spectral fits are subject to considerable uncertainty due to lack of knowledge of the true contribution of hydrogen continuum emission. Maximal mass estimates, i.e., assuming no hydrogen, result in 1.5 $M_\\odot$ of Fe and 0.24 $M_\\odot$ of Si. Under the assumption that an equal quantity of hydrogen has been mixed into the ejecta, we estimate 0.8 $M_\\odot$ of Fe and 0.12 $M_\\odot$ of Si.  These characteristics support the view that in \\dem71\\ we see Fe-rich ejecta from a Type Ia SN several thousand years after explosion. ",
        "introduction": "A distinguishing characteristic of Type Ia supernovae (SNe) is the large amount of iron that they are thought to produce mainly through the decay of radioactive Ni$^{56}$. The nuclear decay of this material powers the SN light curve, thereby providing a constraint on the quantity of material ejected in individual SN events. Although variation in the amount and distribution of radioactive ejecta likely accounts for the modest range of peak absolute magnitudes seen in Type Ia SNe, by and large SNe Ia are strikingly homogeneous in their properties. Various observational constraints (e.g., the absence of hydrogen in optical spectra near optical maximum, occurrence in older stellar populations, and the relatively large amount of Ni$^{56}$ produced) have led to the current framework involving carbon deflagration/detonation in a white dwarf driven to the Chandrasekhar limit by accretion. The most likely progenitor system (Branch et al.~1995) is a C-O white dwarf accreting H/He-rich gas from a companion, either from its wind or through Roche lobe overflow. Remnants of SN~Ia should contain about one Chandrasekhar mass of ejecta comprised of 0.6--0.8 $M_\\odot$ of iron with intermediate-mass elements (O, Mg, Si, S, Ca) distributed in the outer layers (e.g., Iwamoto et al.~1999).  Current models predict that SNe Ia should leave no compact remnant.  One of the goals of supernova remnant (SNR) research has been to use remnant properties to infer the nature of the progenitor. Such properties as the partially neutral nature of the surrounding ambient medium (Tuohy et al.~1982) or the Fe-rich composition of the ejecta (Hughes et al.~1995) have been used to argue for a SN~Ia origin for some SNRs.  In this article we investigate the Large Magellanic Cloud (LMC) SNR \\dem71\\ (0505$-$67.9) which exhibits both of these features. \\dem71\\ (Davies, Elliott, \\& Meaburn 1976) was first proposed as a supernova remnant based on the \\einstein\\ X-ray survey of the LMC (Long, Helfand, \\& Grabelsky 1981).  Follow-up optical spectroscopy confirmed this identification and revealed the remnant to be Balmer-dominated (Tuohy et al.~1982; Smith et al.~1991), that is to say, one whose optical spectrum is dominated by hydrogen emission with little or no emission from forbidden lines of [\\ion{O}{3}] or [\\ion{S}{2}].  The remnants of Tycho's SN and SN1006 are Galactic examples of such SNRs (Kirshner \\& Chevalier 1978; Schweizer \\& Lasker 1978).  \\asca\\ X-ray studies of \\dem71\\ (Hughes, Hayashi, \\& Koyama 1998) yielded an age of $\\sim$5000 yr and evidence for the presence of SN ejecta in the form of an enhanced global abundance of Fe.  Here we present results from new \\chandra\\ observations of the remnant that bear closely on the nature of the originating SN. ",
        "conclusions": "The spatial and spectral analysis of the \\chandra\\ ACIS data reveal \\dem71\\ to be a textbook example of a supernova remnant: an outer blast wave shock propagating through the ISM that surrounds a core of SN ejecta heated to X-ray--emitting temperatures by the inward moving reverse shock. Our analysis strongly argues for a Type Ia SN origin for \\dem71.  Specifically we find (1) a large Fe to O ratio in the ejecta X-ray emission; (2) $\\sim$1.5 $M_\\odot$ of ejecta from dynamical considerations; (3) a relatively low mass ratio of Si to Fe ($\\sim$0.15); and (4) a large mass of ejected Fe ($\\sim$0.8 $M_\\odot$). These results are all consistent with SN~Ia ejecta and further may have some power to discriminate between models for how the flame front propagates in these explosions (see, e.g., Iwamoto et al.~1999).  The ejecta are mostly contained within a thick shell and are partially stratified. According to our deprojection analysis, Fe extends throughout the entire ejecta even into the center, while Si appears to be contained only within the outer half of the ejecta.  A considerably deeper \\chandra\\ observation of \\dem71\\ has been approved for observation in cycle 4.  This will allow us to search for and study spectral variations within the Fe-rich ejecta and to determine its morphology in greater detail. Measuring the kinematics of the Fe-rich ejecta is an important next step. Unfortunately this is likely to be beyond the capabilities of \\chandra, \\xmm, and even \\astroe2, although the Constellation-X Facility, if it attains angular resolution of 10$^{\\prime\\prime}$, should be able to detect the expected motions of $\\sim$100 km s$^{-1}$. We are also pursuing ground-based optical searches for coronal [\\ion{Fe}{14}] $\\lambda$5303 \\AA\\ line emission from the ejecta, which could provide complementary insights into the nature and dynamics of the ejecta."
    },
    "0212/astro-ph0212432_arXiv.txt": {
        "abstract": "{ We present in this paper a substructure and spectroimaging study of the Coma cluster of galaxies based on \\xmm data. \\xmm performed a mosaic of observations of Coma to ensure a large coverage of the cluster. We add the different pointings together and fit elliptical beta-models to the data. We subtract the cluster models from the data and look for residuals, which can be interpreted as substructure. We find several significant structures: the well-known subgroup connected to NGC~4839 in the South-West of the cluster, and another substructure located between NGC~4839 and the centre of the Coma cluster. Constructing a hardness ratio image, which can be used as a temperature map we see that in front of this new structure the temperature is significantly increased (higher or equal 10~keV). We interpret this temperature enhancement as the result of heating as this structure falls onto the Coma cluster. We furthermore reconfirm the filament-like structure South-East of the cluster centre. This region is significantly cooler than the mean cluster temperature. We estimate the temperature of this structure to be equal or below 1~keV. A possible scenario to explain the observed features is stripping caused by the infall of a small group of galaxies located around the two galaxies NGC~4921 and NGC~4911 into the Coma cluster with a non-zero impact parameter. We also see significant X-ray depressions North and South-East of NGC4921, which might either be linked to tidal forces due to the merger with the Western structure or connected to an older cluster merger. ",
        "introduction": "The Coma cluster, also known as Abell 1656, is one of the best studied clusters in our Universe (see Biviano \\cite{biviano} for a historical review). It has been explored in all wavelengths from radio to hard X-rays. However, despite the large number of observations, the physics of the Coma cluster and its dynamical state is not yet completely understood. The cluster hosts a powerful radio halo (Feretti \\& Giovannini \\cite{feretti}) and Beppo-Sax data recently revealed the existence of hard non-thermal X-ray emission (Fusco-Femiano et al. \\cite{fusco-femiano}).  Extensive work has been carried out to measure the velocity dispersion and distribution of the galaxies located in the Coma cluster. Colless \\& Dunn (\\cite{colless}) for example published a catalog of 552 redshifts and determined the mean redshift of the cluster to be $cz = 6853$~km/sec and a velocity dispersion of $\\sigma_{cz} = 1082$~km/sec. Colless \\& Dunn (\\cite{colless}) find and confirm the presence of substructure in the galaxy distribution: there is a sub-group located around NGC~4839, which lies to the South-West of the cluster. The existence of this substructure was originally found by Mellier et al. (\\cite{mellier}) and Merritt \\& Trimbley (\\cite{merritt}). Furthermore, it was noted by Fitchett \\&  Webster (\\cite{fitchett}) that the two central galaxies, NGC~4889 and NGC~4874 have a significant difference in velocity, which is indicative of a recent merger. It is probable that NGC~4889 belonged to a subgroup, which recently merged with the main cluster and that NGC~4874 is the central dominant galaxy of the main cluster.  There are a number of papers addressing the thermal X-ray emission of the Coma cluster which comes from the hot thermal intracluster medium (ICM). White, Briel \\& Henry (\\cite{white}) discussed substructures in the Coma cluster based on ROSAT-data. Briel, Henry \\& B\\\"ohringer (\\cite{briel92}) analyzed the ROSAT All-Sky-Survey data and determined the X-ray surface brightness profile of the cluster out to a radius of roughly 100 arcmin. Hughes (\\cite{hughes}) performed an extensive mass analysis on the Coma cluster based on X-ray and optical data. Later work by Vikhlinin, Forman \\& Jones (\\cite{vikhlinin}) revealed the existence of a filament-like substructure, which is most likely linked to NGC~4911, a bright spiral galaxy located South-East with respect to the centre of the cluster. Spectroimaging studies based on ASCA-data constrained the X-ray temperature distribution of the ICM. Honda et al. (\\cite{vikhlinin}) and Watanabe et al. (\\cite{watanabe}) presented temperature maps of the Coma cluster based on ASCA-data, which clearly showed that the ICM is not isothermal. Furthermore, Donnelly et al. (\\cite{donnelly}) saw indications for a hot spot North of the centre of the Coma cluster. Recently, first results of the Coma cluster based on XMM-Newton data were published: Briel et al. (\\cite{briel01}) discussed the overall morphology of the cluster and showed the existence of several point-like sources, which are linked to individual galaxies in the cluster and which have not been accounted for in previous studies. Arnaud et al. (\\cite{arnaud}) presented a temperature map of the central part of the cluster based on \\xmm data with unprecedented spatial resolution. This map showed again variations of the X-ray temperature (see Jansen et al. \\cite{jansen} for an overview of \\xmm and Turner et al. \\cite{turner} and Strueder et al. \\cite{strueder} for an overview of the EPIC detectors). However, the previously found hot spot found by Donnelly et al. \\cite{donnelly} was not confirmed. Lastly, Neumann et al. (\\cite{neumann}) presented strong evidence based on \\xmm data that the substructure located around NGC~4839 (see above) is on its first infall onto the cluster from the South-West and has not yet passed its centre.  This is in contradiction with previous work by Burns et al. \\cite{burns}, who suggested that the group around NGC~4839 has already passed the centre of the Coma cluster once.  The aim of this paper is to perform a substructure and spectroimaging study of the Coma cluster based on \\xmm data, to better understand the underlying physics. The work presented here is our first attempt to discuss the overall dynamics of the cluster, and is by no means exhaustive. It aims to show the large potential provided by the data obtained with \\xmm . The paper is organized as follows: after the introduction we present the observations in Sec.2. This is followed by the description of our data treatment in Sec.3. In Sec.4 we present our spatial analysis and in Sec.5 we describe our spectro-imaging results. In Sec.6 we discuss our results and finish in Sec.7 with the Conclusion.  Throughout this paper, we assume $H_0=50$~km s$^{-1}$ Mpc$^{-1}$, $\\Omega=1$ and $q_0 = 1/2$. At a redshift of $z=0.0232$, $1^\\prime$ corresponds to 38.9~kpc. The cluster is sufficiently close so that variations of the geometry of the universe do not matter. ",
        "conclusions": "We performed a substructure and spectroimaging study based on the \\xmm data of the Coma cluster. The data show the existence of substructures which are very likely interacting with the main cluster. The combination of spatial and spectral information based on data obtained with \\xmm  allows us to put global constraints on the dynamical state of the Coma cluster. These are important for our understanding of the physics of clusters. The results presented here show the large potential of future studies of clusters of galaxies observed with \\xmm."
    },
    "0212/astro-ph0212375_arXiv.txt": {
        "abstract": "We present a detailed analysis of the two-point correlation function, $\\xisp$, from the 2dF Galaxy Redshift Survey (2dFGRS). The large size of the catalogue, which contains $\\sim 220\\,000$ redshifts, allows us to make high precision measurements of various properties of the galaxy clustering pattern. The effective redshift at which our estimates are made is $z_s \\approx 0.15$, and similarly the effective luminosity, $L_s \\approx 1.4L^{\\ast}$. We estimate the redshift-space correlation function, $\\xis$, from which we measure the redshift-space clustering length, $s_0 = \\as0\\hmpc$. We also estimate the projected correlation function, $\\Xi(\\sigma)$, and the real-space correlation function, $\\xir$, which can be fit by a power-law $(r/r_0)^{-\\gamma_r}$, with $r_0 = \\ar0\\hmpc$, $\\gamma_r = \\agamr$. For $r \\gtrsim 20\\hmpc$, $\\xi$ drops below a power-law as, for instance, is expected in the popular $\\lcdm$ model. The ratio of amplitudes of the real and redshift-space correlation functions on scales of $8 - 30\\hmpc$ gives an estimate of the redshift-space distortion parameter $\\beta$. The quadrupole moment of $\\xisp$ on scales $30 - 40\\hmpc$ provides another estimate of $\\beta$. We also estimate the distribution function of pairwise peculiar velocities, $f(v)$, including rigorously the significant effect due to the infall velocities, and find that the distribution is well fit by an exponential form. The accuracy of our $\\xisp$ measurement is sufficient to constrain a model, which simultaneously fits the shape and amplitude of $\\xir$ and the two redshift-space distortion effects parameterized by $\\beta$ and velocity dispersion, $a$. We find $\\beta = \\betaoverall$ and $a = \\aoverall\\kms$, though the best fit values are strongly correlated. We measure the variation of the peculiar velocity dispersion with projected separation, $a(\\sigma)$, and find that the shape is consistent with models and simulations. This is the first time that $\\beta$ and $f(v)$ have been estimated from a self-consistent model of galaxy velocities. Using the constraints on bias from recent estimates, and taking account of redshift evolution, we conclude that $\\beta(L = L^*, z = 0) = 0.47 \\pm 0.08$, and that the present day matter density of the Universe, $\\om \\approx 0.3$, consistent with other 2dFGRS estimates and independent analyses. ",
        "introduction": "The galaxy two-point correlation function, $\\xi$, is a fundamental statistic of the galaxy distribution, and is relatively straightforward to calculate from observational data. Since the clustering of galaxies is determined by the initial mass fluctuations and their evolution, measurements of $\\xi$ set constraints on the initial mass fluctuations and their evolution.  The astrophysics of galaxy formation introduces uncertainties, but there is now good evidence that galaxies do trace the underlying mass distribution on large scales.  In this paper we analyse the distribution of $\\sim$ 220\\,000 galaxies in the 2-degree Field Galaxy Redshift Survey (2dFGRS, Colless \\etal~2001). A brief summary of the data is presented in Section 2. Much of our error analysis makes use of mock galaxy catalogues generated from $N$-body simulations which are also discussed in Section 2.  The two-dimensional measurement $\\xisp$, where $\\sigma$ is the pair separation perpendicular to the line-of-sight and $\\pi$ is the pair separation parallel to the line-of-sight provides information about the real-space correlation function, the small-scale velocity distribution, and the systematic gravitational infall into overdense regions. The spherical average of $\\xisp$ gives an estimate of the redshift-space correlation function, $\\xis$, where $s=\\sqrt{(\\pi^2 + \\sigma^2)}$, so the galaxy separations are calculated assuming that redshift gives a direct measure of distance, ignoring the effects of peculiar velocities. Integrating $\\xisp$ along the line-of-sight sums over any peculiar velocity distributions, and so is unaffected by any redshift-space effects. The resulting projected correlation function, $\\Xi(\\sigma)$, is directly related to the real-space correlation function.  Our estimates of $\\xisp$, $\\xis$, $\\Xi(\\sigma)$ and $\\xir$ are presented in Section 3. These statistics have been measured from many smaller redshift surveys (\\eg Davis \\& Peebles 1983; Loveday \\etal~1992; Jing, Mo \\& B\\\"orner~1998; Hawkins \\etal~2001; Zehavi \\etal~2002), but since they sample smaller volumes, there is a large cosmic variance on the results. The large volume sampled by the 2dFGRS leads to significantly more reliable estimates. A preliminary analysis was performed on the 2dFGRS by Peacock \\etal~(2001) but we now have a far more uniform sample and twice as many galaxies. Madgwick \\etal~(2003) have measured these statistics for spectral-type sub-samples of the 2dFGRS.  Peculiar velocities of galaxies lead to systematic differences between redshift-space and real-space measurements, and we can consider the effects in terms of a combination of large-scale coherent flows induced by the gravity of large-scale structures, and a small-scale random peculiar velocity of each galaxy (e.g. Marzke \\etal~1995; Jing \\etal~1998). The large-scale flows compress the contours of $\\xisp$ along the $\\pi$ direction, as described by Kaiser (1987) and Hamilton (1992). The amplitude of the distortion depends on the mean density of the universe, $\\om$, and on how the mass distribution is clustered relative to galaxies, which can be parameterized in terms of a linear bias $b$, defined so that $\\delta_g = b\\delta_m$, where $\\delta$ represents fluctuations in the density field. The random component of peculiar velocity for each galaxy means that the observed $\\xisp$ is convolved in the $\\pi$ coordinate with the pairwise distribution of random velocities. In section 4 we describe the construction of a model $\\xisp$ from these assumptions about redshift-space distortions and also the shape of the correlation function.  In section 5 we use the $Q$ statistic (Hamilton 1992) based on the quadrupole moment of $\\xisp$ to estimate the parameter $\\beta\\approx\\om^{0.6}/b$. In the absence of the small-scale random velocities the shape of $\\xisp$ contours on large scales is directly related to the parameter $\\beta$. A similar estimate of $\\beta$ is provided by the ratio of amplitudes of $\\xis$ to $\\xir$ and this is also presented in Section 5.  In section 6, we use the Landy, Szalay \\& Broadhurst (1998) method to estimate the distribution of peculiar velocities. This technique ignores the effect of large-scale distortions and uses the Fourier transform of $\\xisp$ to estimate the distribution of peculiar velocities, $f(v)$.  The large sample volume of the 2dFGRS makes our measurements more reliable than previous estimates in the same way as for the correlation functions mentioned earlier.  These two approaches provide reasonable estimates of $\\beta$ and $f(v)$ so long as the distortions at small and large scales are completely decoupled. This is not the case for real data, and so we have fitted models which simultaneously include the effects of both $\\beta$ and $f(v)$. The resulting best-fit parameters are the most self consistent estimates. Previous data-sets have lacked the signal-to-noise to allow a reliable multi-parameter fit in this way. Our fitting procedure and results are described in Section 7.  In Section 8, we examine the luminosity and redshift dependence of $\\beta$ and combine our results with estimates of $b$ (Verde \\etal~2002; Lahav \\etal~2002) to estimate $\\om$, and compare this with other recent analyses.  In Section 9, we summarise our main conclusions. When converting from redshift to distance we assume the Universe has a flat geometry with $\\oml = 0.7$, $\\om = 0.3$ and $H_0 = 100\\,h\\kmsmpc$, so that all scales are in units of $\\hmpc$. ",
        "conclusions": "In this paper we have measured the correlation function, and various related quantities using 2dFGRS galaxies. Our main results are summarised as follows:  \\renewcommand{\\theenumi}{(\\roman{enumi})} \\renewcommand{\\labelenumi}{\\theenumi}  \\begin{enumerate} \\item{The spherical average of $\\xisp$ gives the redshift-space correlation function, $\\xis$, from which we measure the redshift space clustering length, $s_0 = \\as0\\hmpc$. At large and small scales $\\xis$ drops below a power law as expected, for instance, in the $\\lcdm$ model.}  \\item{The projection of $\\xisp$ along the $\\pi$ axis gives an estimate of the real-space correlation function, $\\xir$, which on scales $0.1 < r < 12\\hmpc$ can be fit by a power-law $(r/r_0)^{-\\gamma_r}$ with $r_0 = \\ar0\\hmpc$, $\\gamma_r = \\agamr$. At large scales, $\\xir$ drops below a power-law as expected, for instance, in the $\\lcdm$ model.}  \\item{The ratio of real and redshift-space correlation functions on scales of $8 - 30\\hmpc$ reflects systematic infall velocities and leads to an estimate of $\\beta = \\betaxirat$. The quadrupole moment of $\\xisp$ on large scales gives $\\beta = \\betaQ$.}  \\item{Comparing the projections of $\\xisp$ along the $\\pi$ and $\\sigma$ axes gives an estimate of the distribution of random pairwise peculiar velocities, $f(v)$. We find that large-scale infall velocities affect the measurement of the distribution significantly and cannot be neglected. Using $\\beta=0.49$, we find that $f(v)$ is well fit by an exponential with pairwise velocity dispersion, $a = 570 \\pm 25\\kms$, at small $\\sigma$.}  \\item{A multi-parameter fit to $\\xisp$ simultaneously constrains the shape and amplitude of $\\xir$ and both the velocity distortion effects parameterized by $\\beta$ and $a$. We find $\\beta = \\betaoverall$ and $a = \\aoverall\\kms$, using the Hubble Volume $\\xir$ as input to the model. These results apply to galaxies with effective luminosity, $L \\approx 1.4L^{\\ast}$ and at an effective redshift, $z_s \\approx 0.15$. We also find that the best fit values of $\\beta$ and $a$ are strongly correlated.}  \\item{We evolve our value for the infall parameter to the present day and critical luminosity and find $\\beta(L=L^*, z = 0) = 0.47 \\pm 0.08$. Our derived constraints on $\\om$ and $b$ are consistent with a range of other recent analyses.}  \\end{enumerate}  Our results show that the clustering of 2dFGRS galaxies as a whole is well matched by a low density $\\lcdm$ simulation with a non-linear local bias scheme based on the smoothed dark-matter density field. Nevertheless, there are features of the galaxy distribution which require more sophisticated models, for example the distribution of pairwise velocities and the dependence of galaxy clustering on luminosity or spectral type. The methods presented have also been used on sub-samples of the 2dFGRS, split by their spectral type (Madgwick \\etal\\ 2003)."
    },
    "0212/astro-ph0212469_arXiv.txt": {
        "abstract": "How massive stars die -- what sort of explosion and remnant each produces -- depends chiefly on the masses of their helium cores and hydrogen envelopes at death.  For single stars, stellar winds are the only means of mass loss, and these are chiefly a function of the metallicity of the star.  We discuss how metallicity, and a simplified prescription for its effect on mass loss, affects the evolution and final fate of massive stars.  We map, as a function of mass and metallicity, where black holes and neutron stars are likely to form and where different types of supernovae are produced.  Integrating over an initial mass function, we derive the relative populations as a function of metallicity.  Provided single stars rotate rapidly enough at death, we speculate upon stellar populations that might produce gamma-ray bursts and jet-driven supernovae. ",
        "introduction": "The fate of a massive star is governed chiefly by its mass and composition at birth and by the history of its mass loss.  For single stars, mass loss occurs as a result of stellar winds for which there exist semi-empirical estimates. Thus, within currently existing paradigms for the explosion, the fate of a star of given initial mass and composition is determined (the Russell-Vogt theorem). If so, one can calculate realization frequencies for stellar explosions and remnants of various kinds and estimate how these might have evolved with time.  Such estimates are fraught with uncertainty. The litany of complications is long and requires discussion (\\Sect{uncer}). No two groups presently agree, in detail, on the final evolution of {\\sl any} massive star (including its explosion energy, remnant mass, and rotation rate) and the scaling of mass loss with metallicity during different evolutionary stages is widely debated. Still it is worthwhile to attempt an approximate table of histories.  We would like to know, within the comparatively well understood domain of stars that do not experience mass exchange with a companion and for a particular set of assumptions regarding mass loss and explosion, what sort of supernova each star produces and what sort of bound remnant, if any, it leaves. If possible, we would also like some indication of which massive stars might make gamma-ray bursts.  In this paper we construct such a table of stellar fates and remnants. In \\Sect{assump} we describe our assumptions regarding mass loss, explosion mechanism(s), and remnant properties, and in \\Sect{uncer} discuss the uncertainties. Section 4 delineates the sorts of stellar explosions and collapses we want to distinguish, and in \\Sect{pop} we discuss the resulting realizations of different outcomes as a function of metallicity in the galaxy. ",
        "conclusions": "\\lSect{con}  We have described, qualitatively, the likely fates of single massive stars as a function of metallicity.  Our results suggest various trends in the observations of these objects which may be subject to observational tests.  \\begin{itemize} \\item Normal Type Ib/c SNe are not produced by single stars until the metallicity is well above solar. Otherwise the helium core mass at death is too large.  This implies that most Type Ib/c SNe are produced in binary systems where the binary companion aids in removing the hydrogen envelope of the collapsing star. \\item Although less extreme than Type Ib/c SNe, single stars also do not produce Type IIL SNe at low metallicities.  Similar to Type Ib/c SNe, Type IIL SNe from single stars are probably ``weak'' SNe until the metallicity exceeds solar, also implying that Type IIL SNe are produced in binaries. \\item If GRBs are produced by single star collapse (perhaps unlikely given the constraints on angular momentum), single stars only make up a small subset of GRB progenitors at higher metallicities.  It is more likely that binary systems form GRBs. Such systems will occur more frequently at low metallicities \\citep{FWH99}. \\item Jet-driven supernovae from single stars are likely to be much more common than GRBs from single stars. \\end{itemize}  It is difficult to make direct comparisons to observations without including binary stars in our analysis, but there are a number of constraints that should be considered.  First, an increasing number of JetSNe and weak supernovae explosions are being discovered \\citep{nak01,SCL98,tur98}.  Although there is an observational bias against the discovery of weak supernovae and they are much dimmer than JetSNe, they may still dominate the sample of stars more massive than 25\\,\\Msun.  Clearly, good statistics (and correct analysis of the systematics) are necessary to determine the relative ratio of jet-driven and weak SNe.  With such statistics, we may be able to place constraints on the rotation of massive stellar cores.  If the IMF becomes more top-heavy at low metallicity ($\\lesssim$$\\,\\Ep{-4}\\,\\Zsun$; \\citealt{BFCL01,sch02}) the number of core collapse supernovae (mostly Type IIp) and GRBs (if occurring in single stars) should significantly increase at high redshift. If the current estimates of a characteristic mass of $\\sim100\\,\\Msun$ for primordial stars \\citep{BCL99,ABN00,NU01} is correct we should expect a large fraction of pair SNe and very massive black holes (or Type III collapsars) at zero metallicity, as well as an increase of massive black holes from stars in the 60--140\\,\\Msun region."
    },
    "0212/astro-ph0212143_arXiv.txt": {
        "abstract": "D Whysong and I have been carrying out a mid-IR survey of 3C radio galaxies to see which have IR-excesses, interpretable as hidden Type-1 AGN.  Mostly due to relentless bad weather, we have a modest collections of observations so far. However, we can say with confidence that some radio galaxies are \\textit{true kinetic, nonthermal emitters}, with negligible energetic contribution from a visible or hidden Big Blue Bump  To make up for the modest collection of mid-IR observations so far, we also present a new spectropolarimetric discovery regarding the underlying shape of the quasar Big Blue Bump in the Balmer Edge region, a fairly radical view of quasar reddening, and a bit of Keck adaptive optics data on Cygnus A and the ultraluminous infrared galaxy NGC6240. ",
        "introduction": "My topic was to be Thermal Emission as a Test for Hidden \\textit{AGN} in Radio Galaxies, work in collaboration with, and in fact largely carried out by UCSB graduate student Dave Whysong.  As I will explain, we are attempting to determine which FRII radio galaxies in the 3C catalog are likely to have hidden quasars by the consequent reradiation of energy by warm dust.  The answer is far from all, and the dependence on optical spectrum, radio size, etc is essential for the Unified Model, and for black hole accretion mechanisms and demographic studies.  Since our data set is insufficiently complete, I will emphasize instead some nearby individual objects of interest.  These latter are discussed in astro-ph/0207385, a manuscript submitted to ApJ.  Four partially baked (about half actually) results together must suffice for one fully baked one.  Therefore I've added some information on a recent discovery by Makoto Kishimoto, Omer Blaes and myself, which reveals the spectral shape of the Big Blue Bump for the first time in the important region of the Balmer edge.  Next I show some results on the quasar reddening curve, which are startling, and very important, or else wrong.  The authors there are Martin Gaskell, Rene Goosmann, David Whysong, and myself.  Finally I preview some adaptive optics work from Keck on Cygnus A and \\textit{NGC6240}, courtesy of collaborators Claire Max, Gabriella Canalizo, David Whysong, Bruce Macintosh, and Alan Stockton.  This paper is obviously not a review but just a compact story that cites a negligible fraction of the relevant literature. ",
        "conclusions": ""
    },
    "0212/astro-ph0212233_arXiv.txt": {
        "abstract": "In an effort to probe the extent of variations in the interstellar $^7$Li/$^6$Li ratio seen previously, ultra-high-resolution (R $\\sim$ 360,000), high signal-to-noise spectra of stars in the Perseus OB2 and Scorpius OB2 Associations were obtained.  These measurements confirm our earlier findings of an interstellar $^7$Li/$^6$Li ratio of about 2 toward $o$~Per, the value predicted from models of Galactic cosmic ray spallation reactions. Observations of other nearby stars yield limits consistent with the isotopic ratio $\\sim$ 12 seen in carbonaceous chondrite meteorites.  If this ratio originally represented the gas toward $o$~Per,  then to decrease the original isotope ratio to its current value an order of magnitude increase in the Li abundance is expected, but is not seen.  The elemental K/Li ratio is not unusual, although Li and K are formed via different nucleosynthetic pathways.  Several proposals to account for the low $^7$Li/$^6$Li ratio were considered, but none seems satisfactory.  Analysis of the Li and K abundances from our survey highlighted two sight lines where depletion effects are prevalent.  There is evidence for enhanced depletion toward X~Per, since both abundances are lower by a factor of 4 when compared to other sight lines.  Moreover, a smaller Li/H abundance is observed toward 20~Aql, but the K/H abundance is normal, suggesting enhanced Li depletion (relative to K) in this direction.  Our results suggest that the $^7$Li/$^6$Li ratio has not changed significantly during the last 4.5 billion years and that a ratio $\\sim$ 12 represents most gas in the solar neighborhood.  In addition, there appears to be a constant stellar contribution of $^7$Li, indicating that one or two processes dominate its production in the Galaxy. ",
        "introduction": "One of the primary goals in astronomy centers on the origin of the elements, but a complete picture remains elusive.  Our focus is an improved understanding of light element synthesis through observations of interstellar lithium in diffuse clouds. Since models of Big Bang nucleosynthesis (BBN) yield 10\\% of the current abundance of $^7$Li and negligible amounts of $^6$Li (Suzuki, Yoshi, \\& Beers 2000), much theoretical effort has gone into finding the source for Li. In the 1970's, the work of Reeves and collaborators (Reeves, Fowler, \\& Hoyle 1970; Meneguzzi, Audouze, \\& Reeves 1971, hereafter (MAR); Reeves et al. 1973; Reeves 1974) provided an alternate method of light element production which plays an important role throughout the history of the Galaxy. This process is now known as standard Galactic cosmic ray (GCR) spallation. Light elements are created when interstellar C, N, and O nuclei are broken apart by Galactic cosmic rays, the spallation process, and by $\\alpha$~$-$~$\\alpha$ fusion reactions.  The exact nature of the cosmic ray source is still under debate (e.g., Ramaty et al. 2000a; Fields et al. 2001). Two commonly accepted creation mechanisms for cosmic rays involve supernovae; p and $\\alpha$ particles are produced in and then accelerated by the explosion or are sputtered off interstellar grains and then accelerated.  A low energy cosmic ray (LECR) component, accelerated C and O nuclei, may also be present (Ramaty, Kozlovsky, \\& Lingenfelter 1996).  Lithium, beryllium, and boron are produced through the inflight fragmentation of C and O nuclei during collisions with ambient interstellar H and He (Cass\\'{e}, Lehoucq, Vangioni-Flam 1995; Higdon, Lingenfelter, \\& Ramaty 1998; Lemoine, Vangioni-Flam, \\& Cass\\'{e} 1998; Lingenfelter, Ramaty, \\& Kozlovsky 1998; Vangion-Flam et al. 1998; Parizot \\& Drury 1999; Ramaty \\& Lingenfelter 1999; Vangioni-Flam, Cass\\'{e}, \\& Audouze 2000).  The inverse of standard GCR reactions, this LECR component originates from core collapse supernovae (SN~{\\small II}) grouped in a superbubble (Higdon et al. 1998; Lingenfelter et al. 1998).  Both standard GCR and LECR/superbubble models are able to account for the present day abundance of $^6$Li, $^9$Be, and $^{10}$B, but only 10 $-$ 25\\% of $^7$Li and about 50\\% of $^{11}$B (e.g. MAR; Ramaty et al. 1997).  The rapid rise in the $^7$Li abundance for stars with [Fe/H]~$\\geq$~-0.5 indicates the existence of a stellar source of $^7$Li.  Proposed stellar sources include asymptotic giant branch (AGB; Smith \\& Lambert 1989, 1990; Plez, Smith, \\& Lambert 1993) and red giant branch (RGB; Smith et al. 1995) stars, with contributions from SN~{\\small II} and possibly novae.  In AGB and RGB stars rapid transport of $^7$Be (Cameron \\& Fowler 1971; Boothroyd, Sackmann, \\& Wasserburg 1994, 1995; Wasserburg, Boothroyd, Sackmann 1995) between the He~burning~shell and the envelope yields $^7$Li. In Type II supernovae (SN II), the flux of neutrinos becomes so great that the spallation of heavy nuclei, C, N, and O, into light elements can occur.  Through spallation in the He and C shells, SN II can produce observable amounts of both $^7$Li and $^{11}$B (Woosley et al. 1990; Woosley \\& Weaver 1995). Another possible thermonuclear source of $^7$Li is novae.  Novae are predicted to play a significant part in the $^7$Li abundance (Starrfield et al. 1978; Romano et al. 2001). The nature of the processes involved in the production of $^7$Li is still unclear.  A goal of our survey of the interstellar $^7$Li/$^6$Li ratio is to shed light on the importance of the various processes.  Interstellar Li was first detected in the early 1970's (Traub \\& Carleton 1973).  Since then there have been relatively few reported observations of interstellar Li~{\\small I}.   The early detections (Traub \\& Carleton 1973; Vanden Bout et al. 1978;  Snell \\& Vanden Bout 1981; Hobbs 1984; White 1986) showed weak interstellar Li~{\\small I} absorption, barely discernible from the noise, with equivalent widths ($W_{\\lambda}$s) on the order of at most a few m$\\mbox{\\AA}$. These early observations showed that the interstellar Li~{\\small I} abundance was similar along different lines of sight.  However, large uncertainties remained due to low signal-to-noise ratios and uncertain corrections for ionization and depletion onto grains.   With the advent of modern detectors and more sophisticated observing techniques, reliable observations of $^7$Li, and the much weaker $^6$Li line, were possible (Ferlet \\& Dennefeld 1984; Lemoine et al. 1993; Lemoine, Ferlet, \\& Vidal-Madjar 1995; Meyer, Hawkins, \\& Wright 1993; Knauth et al. 2000, Howarth et al. 2002).  The Li isotope ratio is crucial for studies of Galactic chemical evolution (Reeves 1993; Steigman 1993) since it is free from uncertain corrections (e.g., ionization and depletion).  Published $^7$Li/$^6$Li ratios in the local ISM indicate that the ratio may vary from the Solar System value of 12.3 (Anders \\& Grevesse 1989), but the results are far from conclusive.  Four different isotope ratios have been published for the line of sight toward $\\zeta$~Oph: $\\geq$ 25 (Ferlet \\& Dennefeld 1984) and 6.8$^{+ 1.4}_{- 1.7}$ for a single component fit to the data (Meyer et al. 1993), Lemoine et al. (1995) obtained 1.4$^{+1.2}_{-0.5}$ ($\\pm$ 0.6) and 8.6 $\\pm$ 0.8 ($\\pm$ 1.4) and Howarth et al. (2002) reported an average $^7$Li/$^6$Li ratio of 13.2 $\\pm$ 6.3 for their two component fit.  Lemoine et al. (1993) found an isotope ratio of 12.5$^{+4.3}_{-3.4}$ toward $\\rho$~Oph, consistent with the Solar System value.  Knauth et al. (2000) derived isotope ratios for two velocity components toward $o$~Per of 1.7 $\\pm$ 0.3 and 3.6 $\\pm$ 0.6 and an isotope ratio of 10.6~$\\pm$~2.9 for a single component toward $\\zeta$~Per. The latter ratio is marginally consistent with the value of 5.5$^{+1.3}_{-1.1}$ reported by Meyer et al. (1993).  As discussed in \\S 2, an accurate template for the velocity components along a given sight line appears to be required before consistent results emerge.  In our earlier study (Knauth et al. 2000), high-resolution (R $\\sim$ 180,000) spectra were obtained on the stars $o$ and $\\zeta$ Per in the Perseus OB2 Association.  These stars were chosen for their proximity to IC~348, an active star-forming region.  These data showed that the $^7$Li/$^6$Li ratio varies from the Solar System value toward $\\zeta$~Per to a value of approximately 2 toward $o$~Per.  $o$~Per resides closer to IC~348 than does $\\zeta$ Per and has an order of magnitude higher flux of cosmic rays (Federman, Weber, \\& Lambert 1996).  Knauth et al. (2000) suggested that the ratio of 2 was clear evidence for newly synthesized lithium toward $o$~Per resulting from GCR spallation reactions.  These data barely resolve the 2 velocity components separated by $\\sim$ 3 km s$^{-1}$ toward $o$~Per.  To refine these findings and to probe directions with interstellar clouds separated by $\\sim$ 1 km s$^{-1}$, we obtained ultra-high-resolution (UHR) observations toward other stars in the Perseus OB2 and Scorpius OB2 Associations.  The stars in Sco OB2 lie close to an active star-forming region approximately centered on $\\rho$~Oph.  Furthermore, $\\lambda$~Ori was observed because it resides in a region of active star formation.  In addition, an UHR survey of interstellar K~{\\small I} (Welty \\& Hobbs 2001) indicated that Li~{\\small I} could possibly be detected toward $\\lambda$~Ori.  20~Aql lies relatively far from active regions of star formation and serves as a useful comparison.  The main goals of the present study were to determine reliable $^7$Li/$^6$Li isotope ratios in order to probe the extent of variations in its value and to help constrain the roots of Li production.  Since $^6$Li is produced primarily through GCR spallation reactions, the GCR contribution to $^7$Li can be accurately determined.  Assuming that the BBN contribution is known from observations of halo stars, the remaining $^7$Li comes from Population {\\small I} stars.  These assumptions enabled interstellar constraints to be placed on the stellar production of $^7$Li. The organization of this paper is as follows.  In \\S2 we describe the observations and the data reduction. The profile syntheses for all data sets are discussed and the results are tabulated in \\S3.  The isotope ratios, Li~{\\small I} and K~{\\small I} column densities, elemental abundances, depletion, and elemental K/Li ratios for each sight line are given in \\S4.  In \\S5 we discuss implications for future Li studies and thermal vs. turbulent broadening.  The results of the $^7$Li/$^6$Li isotope ratios are discussed in \\S6 and \\S7.  In \\S8 constraints on stellar sources of $^7$Li are presented.  Finally, the results are summarized and suggestions for future work are contained in \\S 9. ",
        "conclusions": "The interstellar Li abundance and the $^7$Li/$^6$Li isotope ratio were measured toward several bright stars in the solar neighborhood.  These observations form the highest resolution (R $\\sim$ 360,000) survey of interstellar Li to date. From these data it appears that the Li abundance has not significantly changed since the formation of the Solar System and that the Solar System value of $\\sim$ 12 for the $^7$Li/$^6$Li ratio represents most of the gas in the solar neighborhood.  One exception is the LOS toward $o$~Per.  This sight line exhibits a low $^7$Li/$^6$Li ratio of about 2, the value expected from models of GCR spallation reactions, which is consistent with our previous determination (Knauth et al. 2000) from lower resolution observations.  If a value of about 12 was the starting condition for the gas toward $o$~Per, the decrease in the ratio to its present value requires an increase in the Li abundance by an order of magnitude.  However, an enhancement in the total Li abundance is not seen.  The elemental K/Li ratio for this gas is not unusual either, although K and Li are produced via different nucleosynthetic pathways.  The most likely scenario for the ISM toward $o$~Per is that the initial Li was enhanced through GCR spallation reactions occurring as a result of recent supernovae grouped in a superbubble.  The cloud containing a low isotope ratio contains about 20\\% of the total Li content of the clouds.  Therefore, determination of a total Li abundance may mask the true Li abundance for each interstellar cloud. Although the supernova scenario explains many of the observations, this hypothesis does not address the similarity of the K/Li ratio with results for other sight lines. Fractionation cannot account for the low $^7$Li/$^6$Li ratio because the abundance of LiH is too small.  There is also evidence for enhanced x-ray activity in the cluster IC~348 caused by stellar flares in T Tauri stars (Preibisch et al. 1996).  While $^6$Li can be produced during stellar flares (e.g., Deliyannis \\& Malaney 1995; Ramaty et al. 2000b), the enhanced flaring activity in OB associations was found to play a negligible role.  More typical $^7$Li/$^6$Li ratios can be obtained through the inclusion of extra velocity components that are detected in Na~{\\small I} absorption at 5895 \\AA\\ but not seen in K~{\\small I} absorption at 7699 $\\mbox{\\AA}$, although this comes at the expense of highly disparate Na~{\\small I}/Li~{\\small I} and K~{\\small I}/Li~{\\small I} ratios.  Therefore, the exceptional isotopic ratio toward $o$~Per remains a puzzle. Further observations of other stars in or near IC~348 are needed to probe the extent of the region containing a low $^7$Li/$^6$Li ratio.  Observations were made of X~Per, a nearby x-ray binary.  It is believed that X~Per and other x-ray binary systems are formed after a supernova explosion (e.g., Delgado-Mart\\'{i} et al. 2001).  The direction toward X~Per shows almost an order of magnitude smaller Li and K abundances than toward other sight lines.  These smaller abundances suggest enhanced depletion along the LOS.    Unfortunately, the SNR of the X~Per spectrum was too low to allow a useful probe of the region containing the low isotope ratio.  The stars in Sco OB2 were studied since they closely matched the situation present toward $o$~Per.  These stars are in an OB association where star formation is occurring in the $\\rho$ Oph molecular cloud.  There is evidence of a past supernova because $\\zeta$ Oph is a runaway star (Hoogerwerf et al. 2001).  Why is there no evidence for enhanced $^6$Li production in the presence of active star formation ($\\rho$~Oph molecular cloud) and past SN~{\\small II} explosions?  Is the LOS toward $o$~Per really that rare?  More observations are necessary to determine if we are experiencing a selection effect; to date, only three sight lines in each OB association have been studied with sufficient precision.  An additional result of our study involves the LOS toward 20~Aql, which resides far from active star-forming regions.  The $^7$Li/$^6$Li ratio is similar to the Solar System value, but the K/Li ratio is almost double the expected value.  The K/H abundance is similar to the abundance found for other sight lines, indicating that there is less Li. The LOS toward 20~Aql lies on the edge of Radio Loop I (Hayakawa et al. 1979; Sembach, Savage, Tripp 1997), a superbubble. However, the timescale needed to destroy Li through proton burning is significantly longer than the age of bubble.  Another explanation is required to account for the low Li/H abundance.  The low Li abundance toward 20~Aql could be evidence for Li dilution in an old superbubble (Parizot \\& Drury 1999; Knauth et al. 2000).  Observations of Li~{\\small I} and K~{\\small I} toward other stars in the vicinity of 20~Aql are necessary to gain further insight.  We place additional interstellar constraints on the $^7$Li produced by a stellar source.  Our results indicate that there is essentially a constant production of $^7$Li in stars in the solar neighborhood, suggesting one or two processes dominate.  Observations of interstellar and stellar $^7$Li, s-process, and r-process elements can be of great value in constraining the Li production by the various stellar sources.  For future studies of interstellar Li, knowledge of the K~{\\small I} abundance in conjunction with the molecular content can be used to find sight lines that contain observable amounts of Li.  For instance, the results of this survey suggest other targets in Sco OB2.  These stars, which are fainter than X~Per, can be observed at lower resolving power (R $\\sim$ 180,000) because their interstellar spectra reveal simple velocity structure. Preliminary results for these stars indicate that the $^7$Li/$^6$Li ratios are similar to the Solar System value.  Thus, lines of sight with low $^7$Li/$^6$Li ratios appear to be rare.  Observations of stars in the vicinity of IC~348 (in Per OB2) are planned to help clarify the picture for gas toward $o$~Per.  In addition, data from the $Hubble~Space~Telescope$ on the interstellar $^{11}$B/$^{10}$B ratio toward stars in Per OB2 are being analyzed.  The interstellar abundance of $^9$Be and the $^{11}$B/$^{10}$B ratio (Hebrard et al. 1997; Lambert et al. 1998) are crucial to our understanding of Li evolution. Knowledge of $^9$Be and $^{10}$B further constrain the amount of GCR spallation occurring along the line of sight, since GCR spallation is the only production route for these light elements.  Observations of interstellar $^9$Be are not yet feasible because its gas phase abundance is extremely low ($^9$Be/H $\\leq$ 7 $\\times$ 10$^{-13}$; Hebrard et al. 1997). The amount of $^{11}$B also yields constraints on SN~{\\small II} explosions because about 50\\% of $^{11}$B is thought to arise from neutrino-induced spallation reactions (Woosley et al. 1990; Woosley \\& Weaver 1995).    Only through a more complete study of the abundances for a variety of atomic species will the light element puzzle be resolved."
    },
    "0212/astro-ph0212005_arXiv.txt": {
        "abstract": "LXeGRIT is the first prototype of a novel concept of Compton telescope, based on the complete 3-dimensional reconstruction of the sequence of interactions of individual \\g-rays in one position sensitive detector. This balloon-borne telescope consists of an unshielded time projection chamber with an active volume of 400~cm$^2 \\times 7$~cm filled with high purity liquid xenon. Four VUV PMTs detect the fast xenon scintillation light signal, providing the event trigger. 124 wires and 4 anodes detect the ionization signals, providing the event spatial coordinates and total energy. In the period 1999 -- 2001, LXeGRIT has been extensively tested both in the laboratory and at balloon altitude, and its response in the MeV region has been thoroughly characterized. Here we summarize some of the results on pre-flight calibration, event reconstruction techniques, and performance during a 27 hour balloon flight on October 4 -- 5. We further present briefly the on-going efforts directed to improve the performance of this prototype towards the requirements for a base module of a next-generation Compton telescope. ",
        "introduction": "\\label{sec:TPC_approach} Between the energy regimes of photoabsorption and pair production, about 250~\\keV\\ and 6~\\MeV\\ in xenon, \\g-rays interact in matter predominantly via the Compton process. A Compton telescope is thus the most promising approach to measure efficiently the distribution of cosmic sources emitting in the MeV energy band. In a double scatter telescope such as COMPTEL~\\cite{VSchon:93}, the restriction on the event topology to two interactions recorded in two separate detector planes, results in a relatively low detection efficiency ($\\ll 1\\%$).  This figure is greatly improved in a homogeneous, self-triggered, three-dimensional (3D) position sensitive detector such as a liquid xenon time projection chamber (LXeTPC), as proposed by Aprile~et~al.\\cite{EAprile:89:ITNS}. In a TPC of large sensitive volume, passive materials are minimized and a variety of different event topologies are used for imaging, substantially increasing detection efficiency. With the 3D position sensitivity, single-site events are easily identified and rejected as background, as they are not useful for imaging. Multiple-site events are associated with MeV \\g-rays from both source and background. As for any Compton telescope, the imaging capability greatly enhances signal over background compared to formerly employed simple photon counters. In addition, the wealth of information recorded allows for more sophisticated event selections based on Compton kinematics, interaction locations, or interaction separation.    In absence of a time-of-flight measurement, the correct order of the multiple interactions has to be determined from the redundant kinematical and geometrical information measured for each event. The efficiency of this sequence reconstruction is directly determined by the accuracy of the energy and position measurements. As previously shown in Aprile~et~al.\\cite{EAprile:93:monte_carlo} and more recently by other groups \\cite{GJSchmid:99:tracking,SBoggs:00:comp_reconst,UOberlack:00:spie00:CSR}~, event reconstruction based on Compton kinematics allows to correctly order the interactions and to discriminate against background, e.g., from non-Compton sequences or multiple coincident gamma-rays.    The LXeGRIT is the outcome of a systematic program initiated at Columbia University with NASA support to develop the LXeTPC technology for MeV astrophysics.  Experiments with LXe detectors were carried out to measure charge and light yields, energy and position resolution, electron drift velocity and mobility, etc. \\cite{EAprile:91:performance}~.  Following these basic studies, a LXeTPC, with a sensitive volume of 2800 cm$^3$, was developed and its performance established with a varied of \\g-ray sources. The same TPC is used for the balloon-borne LXeGRIT instrument.  The detector works as follows. Both ionization and scintillation signals are detected to measure the event spatial coordinates and the energy. The fast UV (175~nm) Xe light is detected by four PMTs viewing the sensitive volume from below, through quartz windows. The OR of the PMT signals is the TPC trigger signal, marking time zero of each event.  The charge signals induced by free electrons drifting under the applied electric field are detected on two orthogonal planes of parallel wires, providing X-Y position information with millimeter accuracy. The wires pitch is 3~mm, and the separation between X and Y planes is also 3~mm.  Below the wires, at a distance of 3~mm, four anodes are used to collect the total charge liberated in the event. The charge is directly proportional to energy.  The 62 X-wires and 62 Y-wires and the anode are amplified and digitized at 5~MHz, with 8 bit precision for the wires and 10 bit precision for the anodes.  The digitized data are stored with 256 samples per event, covering more than the maximum drift time of 35 $\\mu$s.  The absolute drift time measurement gives the event depth of interaction (Z-position), with an accuracy of 300 $\\mu$m.  This charge readout was designed and tested to image the point-like ionization clouds ($<$~1mm) produced by low energy Compton electrons and photoelectrons which are typical of MeV \\g-ray interactions in the dense LXe. The signal size associated with these clouds is very small (W=15.6~eV compared to W=2.96~eV for Ge and W=3.62~eV in Si). A key prerequisite for this detector to work is to minimize the signal loss due to electron trapping by impurities in the liquid. In LXeGRIT free electrons are drifted over the TPC maximum distance of 7~cm with very little attenuation, and the remaining 5\\% maximum charge loss for full drift length is corrected for based on the known Z-position.  The TPC is enclosed in a cylindrical vessel filled with 7 liters of pure LXe kept at $\\sim -95\\deg$C by a controlled flow of LN$_2$ through a condenser. The vessel is thermally insulated with a vacuum cryostat. The lower section of the cryostat houses the four PMTs of the xenon light readout, as well as the HV distribution circuitry for the wires and the cathode. ",
        "conclusions": ""
    },
    "0212/astro-ph0212555_arXiv.txt": {
        "abstract": "{ In popular cold dark matter cosmological scenarios, stars may have first appeared in significant numbers around a redshift of 10 or so, as the gas within protogalactic halos with virial temperatures $T_\\vir\\gta 10^{4.3}\\,$K (corresponding to masses comparable to those of present--day dwarf ellipticals) cooled rapidly due to atomic processes and fragmented. It is this `second generation' of subgalactic stellar systems, aided perhaps by an early population of accreting black holes in their nuclei, which may have generated the ultraviolet radiation and mechanical energy that ended the cosmic ``dark ages'' and reheated and reionized most of the hydrogen in the universe by a redshift of 6. The detailed history of the universe during and soon after these crucial formative stages depends on the power--spectrum of density fluctuations on small scales and on a complex network of poorly understood `feedback' mechanisms, and is one of the missing link in galaxy formation and evolution studies. The astrophysics of the epoch of first light is recorded in the thermal state, ionization degree, and chemical composition of the intergalactic medium, the main repository of baryons at high redshifts.}  \\addkeyword{cosmology} \\addkeyword{diffuse radiation} \\addkeyword{intergalactic medium}   \\begin{document} ",
        "introduction": "At epochs corresponding to $z\\sim 1000$ the intergalactic medium (IGM) is expected to recombine and remain neutral until sources of radiation and heat develop that are capable of reionizing it. The detection of transmitted flux shortward of the \\lya wavelength in the spectra of sources at $z\\gta 5$ implies that the hydrogen component of this IGM was ionized at even higher redshifts. The increasing thickening of the \\Lya forest recently measured in the spectra of SDSS $z\\sim 6$ quasars (Becker \\etal 2001; Djorgovski \\etal 2001) may be the signature of the trailing edge of the cosmic reionization epoch. It is clear that substantial sources of ultraviolet photons and mechanical energy, like young star--forming galaxies, were already present back then. The reionization of intergalactic hydrogen at $z\\gta 6$ is unlikely to have been accomplished by quasi--stellar sources: the observed dearth of luminous optical and radio--selected QSOs at $z>3$ (Shaver \\etal 1996; Fan \\etal 2001), together with the detection of substantial Lyman--continuum flux in a composite spectrum of Lyman--break galaxies at $\\langle z\\rangle=3.4$ (Steidel et al. 2001), both may support the idea that massive stars in galactic and subgalactic systems -- rather than quasars -- reionized the hydrogen component of the IGM when the universe was less than 5\\% of its current age, and dominate the 1 ryd metagalactic flux at all redshifts greater than 3. An episode of pregalactic star formation may also provide a possible explanation for the widespread existence of heavy elements (like carbon, oxygen, and silicon) in the IGM. There is mounting evidence that the double reionization of helium may have occurred later, at a redshift of 3 or so (see Kriss et al. 2001, and references therein): this is likely due to the integrated radiation emitted above 4 ryd by QSOs.  Establishing what ended the dark ages and when is important for determining the impact of cosmological reionization and reheating on several key cosmological issues, from the role reionization plays in allowing protogalactic objects to cool and make stars, to determining the thermal state of baryons at high redshifts and the small--scale structure in the temperature fluctuations of the cosmic microwave background. Conversely, probing the reionization epoch may provide a means for constraining competing models for the formation of cosmic structures: for example, popular modifications of the CDM paradigm that attempt to improve over CDM by suppressing the primordial power--spectrum on small scales, like warm dark matter (WDM), are known to reduce the number of collapsed halos at high redshifts and make it more difficult to reionize the universe (Barkana \\etal 2001). In this talk I will summarize some recent theoretical developments in understanding the astrophysics of the epoch of first light and the impact that some of the earliest generations of stars, galaxies, and black holes in the universe may have had on the IGM. ",
        "conclusions": ""
    },
    "0212/astro-ph0212413_arXiv.txt": {
        "abstract": "%  We discuss the development of a Java toolbox for astronomical time series data. Rather than using methods conventional in astronomy (e.g., power spectrum and cross-correlation analysis) we employ rule discovery techniques commonly used in analyzing stock-market data. By clustering patterns found within the data, rule discovery allows one to build predictive models, allowing one to forecast when a given event might occur or whether the occurrence of one event will trigger a second. We have tested the toolbox and accompanying display tool on datasets (representing several classes of objects) from the RXTE All Sky Monitor. We use these datasets to illustrate the methods and functionality of the toolbox. We also discuss issues that can come up in data analysis as well as the possible future development of the package. ",
        "introduction": "Many types of variable objects exist in the universe, including stars with predictable behavior (e.g., Cepheids), objects with behavior that is inherently unpredictable (e.g, AGN), and objects with both predictable and irregular variability patterns (e.g., X-ray binaries).  Constant monitoring of variable objects has been a continuing interest in astronomy, beginning with 16th century astronomer David Fabricius, and extending through history to Herschel, Leavitt and others.  Today, monitoring is done by a wide variety of techniques, observers and instruments, from dedicated amateurs, to professional astronomers interested in intensive monitoring of individual objects, to all-sky monitors such as the RXTE ASM and BATSE aboard CGRO.  While they are a new tool, already all-sky monitors have have made important, if not decisive contributions to solving some of astronomy's most persistent mysteries, such as the cosmological origin of gamma-ray bursts and linking emission regions in AGN.  With major initiatives such as the Large-Area Synoptic Survey Telescope (LSST) and Supernova Acceleration Probe (SNAP), all-sky monitors are poised to become a major discovery tool in astronomy. To maximize the utility of large monitoring programs, it is important to devise ways of handling large amounts of data in real time and find not only variability but also predictive patterns among these large data streams. It was with these goals in mind that we undertook this project. ",
        "conclusions": ""
    },
    "0212/astro-ph0212139_arXiv.txt": {
        "abstract": "{We study analytically and numerically the stability of the pressure-less, viscously spreading accretion ring. We show that the ring is unstable to small non-axisymmetric perturbations. To perform the perturbation analysis of the ring we use a stretching transformation of the time coordinate. We find that to 1st order, one-armed spiral structures, and to 2nd order additionally two-armed spiral features may appear. Furthermore, we identify a dispersion relation determining the instability of the ring. The theoretical results are confirmed in several simulations, using two different numerical methods. These computations prove independently the existence of a secular spiral instability driven by viscosity, which evolves into persisting leading and trailing spiral waves. Our results settle the question whether the spiral structures found in earlier simulations of the spreading ring are numerical artifacts or genuine instabilities. ",
        "introduction": "In the theory of accretion discs, the idealised problem of a viscously spreading, pressure-less ring orbiting a central point mass on Keplerian orbits is often used to exemplify the main features of an evolving thin accretion disc, i.e. the inward transport of mass and outward transport of angular momentum \\citep{1981ARA&A..19..137P,FrankKingRaine}. With the approximation of a small kinematic viscosity $\\nu$ that is independent of the surface density, an analytic solution exists for this problem, stated originally by \\citet{lues52} and later by \\citet{1974MNRAS.168..603L}.  This axisymmetric analytic solution of the viscous dust ring is frequently used as test problem for numerical methods developed to simulate accretion discs. The tested algorithms cover particle methods like smoothed particle hydrodynamics (SPH) \\citep[e.g.][]{1994ApJ...431..754F, Spei99} as well as grid-based codes like {\\it RH2D} \\citep[e.g.][]{1999MNRAS.303..696K}.  However, numerical simulations of the evolving ring often show the appearance of additional structures in the disc. In the majority of cases, these structures consist of non-axisymmetric one-armed spirals, but eventually also narrow concentric rings appear. Since their first discovery, it is controversially debated whether these structures are numerical artifacts or mathematical instabilities. While \\citet{1996PASA...13...66M} show that for the axisymmetric case the concentric rings found in SPH simulations are numerical artifacts of the method, we demonstrate in this paper that the non-axisymmetric spiral structures result from a genuine physical instability of the problem.  Recently, \\citet{2001MNRAS.325..231O} analysed radially extended accretion discs in a work presenting a general model of eccentric accretion discs, of which the viscous ring is a simple special case. However, because his main objective was much wider, he did not elaborate on the stability properties of the ring.  The remainder of the paper is organised as follows. In the next section, we first summarise the basic hydrodynamic equations used for our analysis of the evolving disc, and present for reference the analytic solution of the viscously spreading dust ring.  In Sect.~\\ref{sec:perturbation} we use a special perturbation technique, the time-stretching approach, which allows to obtain solutions of the spreading ring at different orders of a small expansion parameter. We derive a linear evolution equation for an eccentricity function $E$. It is shown that to 2nd order, the general solution allows the development of one-armed as well as two-armed spiral features.  In Sect.~\\ref{sec:local} we perform a local stability analysis of the final equation for $E$ and deduce a dispersion relation which indicates an instability for some viscosity laws. These results are compared with numerical simulations we present in Sect.~\\ref{sec:numerics} using two different numerical methods, one grid-based and the other particle based. We find in general very good agreement of both numerical methods with the predictions of the local analysis. Thus, the simulations confirm the existence of a physical instability in the viscously evolving ring problem. Finally, in \\ref{sec:conclusion} we give our conclusions. ",
        "conclusions": "\\label{sec:conclusion} We have shown that the spiral instability found in various simulations of the viscously evolving dust ring can be understood in terms of a secular spiral instability driven by viscosity. To accomplish this, we performed a perturbation analysis using a time-stretching transformation. It turned out that the widely known analytic solutions for the surface density evolution and the azimuthal velocity of the viscously spreading ring consist of the zeroth order terms of the expansion, while the solution for the radial velocity is given by the 1st order terms.  As a result of the perturbation analysis we found that the ring may develop one-armed spiral structures to 1st order and one- and two-armed spiral features to 2nd order. A local stability analysis of an eccentricity function $E$ provided a dispersion relation that led for a constant kinematic viscosity to a growth rate for secular spiral instabilities of \\begin{displaymath} \\sigma_i = \\frac{1}{3} \\, k^2 \\, \\nu \\, + \\, \\mathcal{O}(\\nu/R^2) \\end{displaymath} in the limit $kR \\gg 1$. The exact form of the growth rate depends on the viscosity prescription $\\nu = \\nu (R)$. Our theoretical results could be confirmed in several simulations using two different numerical methods.  As a consequence for numerical simulations of accretion discs, the spiral instability has to be taken into account when using the spreading ring for instance as test problem for code development. Otherwise, occasionally emerging spiral features may be mistaken for numerical instabilities of the algorithms used for the calculation of the viscous forces.  The physical implications of our results are not as obvious. Because the kinematic viscosity $\\nu$ is held axisymmetric throughout the whole stability analysis, the discovered non-axisymmetric instabilities may be of importance mainly in accretion discs where the viscosity is determined uniformly by external effects like irradiation."
    },
    "0212/hep-ph0212126_arXiv.txt": {
        "abstract": "The KamLAND reactor antineutrino experiment has detected a $3.4\\sigma$ flux suppression relative to the expectation if no neutrino oscillations occur. We combine KamLAND data with solar neutrino data and show that the LMA solution is the only viable oscillation solution to the solar neutrino problem at the $4.4\\sigma$ C.~L. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/astro-ph0212407_arXiv.txt": {
        "abstract": "Using a combination of near-infrared and optical photometry, along with multi-object spectroscopy, we have confirmed the existence of a high-redshift cluster of galxies at $z = 0.96$.  The cluster was found using a wide-angle tailed radio source selected from the VLA {\\it FIRST} survey as a cluster signpost.  These types of radio sources are often found in clusters, and are thought to attain their C-shaped morphologies from the relative motion between the radio source host galaxy and the intracluster medium. We present optical/near-infrared color-magnitude diagrams which show a concentration of cluster galaxies in color space.  We also include spectroscopic results obtained from the Keck II LRIS.  Ten galaxies are confirmed at the cluster redshift, with a line-of-sight velocity dispersion of $\\sigma_{\\|}=530^{+190}_{-90}$ km s$^{-1}$, typical of an Abell richness class 0 cluster. Using data from the {\\it ROSAT} public archive, we limit the X-ray luminosity for the cluster to $L_{X,bol} \\la 3 \\times 10^{44}$ erg s$^{-1}$, consistent with the value expected from the $L_{X}-\\sigma$ relation. ",
        "introduction": "Clusters of galaxies trace the largest gravitationally-bound mass concentrations in the Universe; their epoch of formation, and thus their abundance as a function of redshift, is a strong constraint on $\\Omega_M$. High-redshift clusters are very difficult to find -- the faint optical magnitudes of the associated galaxies and the X-ray emission from the intracluster medium (ICM) will often be missed by current surveys at redshifts of $z\\approx1$. Any new clusters revealed at these distances add significantly to the few that are known, and are useful for studies of cosmology and galaxy evolution.  Using a unique method based on radio galaxy morphology, we have confirmed the existence of dozens of clusters with redshifts up to $z=0.85$ (Blanton et al.\\ 2000, 2001). Bent double-lobed radio sources (particularly FR I-type sources [Fanaroff \\& Riley 1974], including wide-angle tails [WATs] and narrow-angle tails [NATs]) are thought to achieve their distinct ``swept-back'', C-shaped morphologies as a result of interaction with dense gas in which they are embedded.  The canonical model describes the radio source's host galaxy as having a significant peculiar velocity relative to the cluster with which it is associated -- the ram pressure of the dense ICM then pushes back the lobes of the radio source giving it a ``bent'' appearance (Owen \\& Rudnick 1976). Often, the host galaxies of WAT radio sources are cD ellipticals, thought to reside at the bottom of cluster potential wells, and are thus not expected to have significant peculiar velocities.  An alternate explanation (Burns et al.\\ 1994), posits bulk motion of the ICM as a result of a recent or ongoing cluster-cluster (or sub-cluster) merger.  The ram pressure of the merging gas then bends the lobes.  Either way, bent-double radio sources are often pointers to clusters, and they may be flags in particular for clusters that have recently formed or merged.  We have selected 384 bent-double (mostly WAT) radio sources from the first 3000 deg$^2$ of the VLA {\\it FIRST} (Faint Images of the Radio Sky at Twenty-cm) survey (Becker, White, \\& Helfand 1995).  Imaging and multi-slit mask spectroscopic follow-up on a moderate-to-high redshift subset of ten of these revealed eight clusters with redshifts as high as $z=0.85$ (Blanton et al.\\ 2000). A complete (area- and magnitude-limited) low-redshift subset including 40 of the 384 objects was observed optically, and $\\sim50\\%$ of them were found to be in richness class Abell 0 or greater clusters (Blanton et al.\\ 2001).  The majority of the low-redshift bent doubles that are found in clusters, based on our optical richness measurements, are detected in the ROSAT All Sky Survey (five out of six detected for $z<0.2$). Those that are not found in clusters, based on our richness measurements, are not detected in the RASS (only one out of nine detected for $z<0.2$).  In other words, the combination of radio and optical data gives us a very clear indication for the presence, or lack of, a cluster, as confirmed with X-ray observations.  In order to identify high-$z$ candidates, we obtained $R$-band images of fields surrounding bent doubles that were blank to the limit of the Digitized Sky Survey (DSS).  Some of these had no optical identifications even down to the limit of our images ($m_{R}\\sim22-23$).  We then performed near-infrared (NIR) observations at the positions of these sources in hopes of uncovering very red, high-$z$ cluster galaxies.  NIR observations have been successful in the past in identifying high-$z$ clusters. Stanford et al.\\ (1997, 2001) found one of the most distant clusters known at $z=1.27$ using NIR and optical imaging, while Rosati et al.\\ (1999) used a combination of X-ray, NIR, and optical observations to confirm the existence of a cluster at $z=1.26$ that may be part of a superstructure including the Stanford et al.\\ (1997) cluster. Using the NIR alone, or in conjunction with radio and/or X-ray data, gives the potential for locating clusters that are at such high-$z$ they would be missed by all but the deepest optical observations. Recently, Gladders \\& Yee (2000) have performed a high-redshift cluster search by using optical and NIR photometry and identifying clusters by their ``red sequence,'' or clustering on a color-magnitude diagram.  Chapman et al.\\ (2000) have used a combination of radio observations with optical and NIR photometry to identify likely clusters at redshifts greater than 0.8.  Here we present optical and NIR observations taken at the position of the bent-double radio source, BD1137+3000.  A combination of imaging and spectroscopy confirms that this source is located in a distant cluster of galaxies with a redshift $z=0.96$.  We assume $H_{\\circ} = 70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M = 0.3$, and $\\Omega_{\\Lambda} = 0.7$ throughout. ",
        "conclusions": "We have provided evidence at optical and near-infrared wavelengths for a high-redshift cluster of galaxies associated with a wide-angle tailed radio galaxy at $z = 0.96$.  Color-magnitude diagrams reveal an over-density of red objects near the radio source with colors consistent with elliptical galaxies at this redshift.  Objects are clustered on the color-magnitude diagrams in the range $4 < (R-K) < 6$ and $1.5 < (J-K) < 2.5$.  Using spectroscopy from the Keck II LRIS, we have confirmed a total of ten galaxies at the cluster redshift.  The line-of-sight velocity dispersion is $\\sigma_{\\|}=530^{+190}_{-90}$ km s$^{-1}$, typical of an Abell richness class 0 cluster.  Using data from the {\\it ROSAT} public archive, we limited the cluster's bolometric X-ray luminosity to $L_{X,bol} \\la 3 \\times 10^{44}$ erg s$^{-1}$, which is consistent with the value expected from the $L_{X}-\\sigma$ relation.  The majority of high-redshift clusters found in traditional X-ray and optical surveys select the most X-ray luminous and optically rich clusters.  The radio selection technique employed here is an important complement to those methods, in part because it can locate high-z clusters with a wide range of X-ray luminosities and optical richnesses, thus aiding in the study of galaxy evolution and cluster formation.  E. L. B. thanks Adam Stanford and Craig Sarazin for helpful discussions regarding the NIR and X-ray data reduction, respectively. Support for E. L. B. was provided by NASA through the {\\it Chandra} Fellowship Program, grant award number PF1-20017, under NASA contract number NAS8-39073. Part of the work reported here was done at the Institute of Geophysics and Planetary Physics, under the auspices of the US Department of Energy by Lawrence Livermore National Laboratory under contract W-7405-Eng-48. Support for M. D. G. was provided by a grant from the National Science Foundation (AST-99-70884). The {\\it FIRST} project is supported by grants from the National Geographic Society, the National Science Foundation (AST 00-98259 and AST 00-98355), NASA (NAG 5-6035), NATO, IGPP, Columbia University, and Sun Microsystems."
    },
    "0212/astro-ph0212541_arXiv.txt": {
        "abstract": "We discuss a concept of off-centred cavity supernova explosion as applied to neutron star/supernova remnant associations and show how this concept could be used to preclude the anti-humane decapitating the Duck (G\\,5.4$-$1.2 $+$ G\\,5.27$-$0.9) and dismembering the Swan (Cygnus Loop), as well as to search for a stellar remnant associated with the supernova remnant RCW\\,86. ",
        "introduction": "Massive ($\\geq 8-10 \\, M_{\\odot}$) stars are the progenitors of most of supernovae (SNe). The explosion of a massive star results in the origin of an extended (tens of parsecs) diffuse supernova remnant (SNR) and a compact stellar remnant. In most cases the stellar remnant is a neutron star (NS). Therefore most of SNRs should be associated with NSs. A NS could be located within the confine of the associated SNR or beyond it, depending on the kick velocity received by the star, the age of the system and some other factors (see below). However, only a small fraction of known SNRs was found to be associated with NSs. Moreover, it is believed that some of proposed NS/SNR associations are the result of geometrical projection. Thus, the obvious lack of associations should be explained or filled up. The latter seems to be more attractive and fruitful in view of the recent splash of discoveries of NSs in SNRs.  The reliability of NS/SNR associations is usually assessed with help of several criteria (see e.g. Kaspi 1996). Some of them are trivial, i.e. should be fulfilled for any association. Other criteria are based on the use of the standard Sedov-Taylor model of SNRs and therefore are over-simplified. The application of these criteria can lead to rejection of genuine associations (Gvaramadze 2000, 2002a,b; Bock \\& Gvaramadze 2002). The point is that the massive stars strongly modify their environs by virtue of their winds and ionizing emission, and it is the subsequent interaction of SN blast waves with their processed ambient medium that results in the observed SNRs (e.g. Shull et al. 1985; Ciotti \\& D'Ercole 1989; Chevalier \\& Liang 1989). It is clear that the presence of circumstellar and interstellar structures could strongly affect the standard sequence and duration of evolutionary stages of the SN blast wave (e.g. Woltjer 1972). For example, the Sedov-Taylor stage could be absent at all if the SN exploded within the wind-blown cavity surrounded by a massive shell (e.g. Franco et al. 1991). Another important point is the proper motion of massive stars, which causes them to explode far from the geometrical centres of their cavities and makes the cavities and other circumstellar and interstellar structures non-spherically-symmetrical. The natural consequence of an off-centred cavity SN explosion is that the SN blast centre does not coincide with the geometrical centre of the future SNR. Taking into account these considerations allows us to enlarge the circle of possible NS/SNR associations and to explain the morphological peculiarities of SNRs (Gvaramadze 2000, 2002a,b; Bock \\& Gvaramadze 2002; Gvaramadze \\& Vikhlinin 2003). On the other hand, the better understanding of origin of peculiar SNRs helps to infer the ``true\" SN explosion sites in these remnants, and therefore to search for new NSs associated with them (Gvaramadze 2002b; Gvaramadze \\& Vikhlinin 2002, 2003; see also Sect.\\,4).  To illustrate the significance of the concept of off-centred cavity SN explosion we show how this concept could be used to preclude the anti-humane decapitating the Duck (G\\,5.4$-$1.2 $+$ G\\,5.27$-$0.9) and dismembering the Swan (Cygnus Loop, G\\,74.0$-$8.5) recently attempted, respectively, by Thorsett, Brisken, \\& Goss (2002) and Uyaniker et al. (2002), and to search for a NS associated with the SNR RCW\\,86 (MSH\\,14-6{\\it3}, G\\,315.4$-$2.30). ",
        "conclusions": ""
    },
    "0212/astro-ph0212294_arXiv.txt": {
        "abstract": "We study the evolution of the relation between half-light (effective) radius, \\re, and mean surface brightness, \\mie, known as Kormendy relation, out to redshift $\\mathrm{z=0.64}$ in V-band restframe on the basis of a large sample of spheroidal galaxies ($\\mathrm{N=228}$) belonging to three clusters of galaxies.  The present sample constitutes the largest data set for which the Kormendy relation is investigated up to a look-back time of $\\mathrm{\\sim 6~Gyr }$ ($ \\rm H_0= 70~Km s^{-1} Mpc^{-1}, \\Omega_M=0.3, \\Omega_{\\Lambda}=0.7$). A new fitting procedure, which suitably accounts for selection criteria effects, allows for the first time to study the trend of the slope ($\\beta$) and of the intrinsic dispersion (\\smie) of the Kormendy relation, and the properties of the whole distribution in the \\re \\, -- \\mie \\, plane as a function of look-back time.  The slope $\\beta$ of the relation does not change from $\\mathrm{z=0.64}$ to the present epoch: $\\beta=2.92\\pm0.08$, implying a tight constraint of $18$--$28\\%$ on the variation of the stellar formation epoch along the sequence of spheroidal galaxies per decade of radius. The intrinsic dispersion of the relation, $\\mathrm{\\sigma^{(i)}_{<\\!\\mu\\!>_e} =0.40\\pm0.03}$, does not vary with redshift and the distribution of galaxy sizes as well as the distribution in the plane of the effective parameters do not vary among the clusters, as proven by the Kolmogorov--Smirnov tests.  We conclude that, whatever the mechanism driving galaxy evolution is, it does not affect significantly the properties of bright galaxies in the \\lre \\, -- \\mie \\, plane at least since $\\mathrm{z=0.64}$.  The evolution of the zeropoint of the Kormendy relation is fully explained by the cosmological dimming in an expanding universe plus the passive luminosity evolution of stellar populations with high formation redshift ($\\mathrm{z_f>2}$). ",
        "introduction": "Early--type galaxies (ETGs) are known to populate a two-dimensional manifold in the space of their observed quantities. This Fundamental Plane (FP) is usually expressed by a correlation among the effective radius \\re, the mean surface brightness \\mie \\, within \\re, and the central velocity dispersion \\s0: \\begin{equation} \\mathrm{ \\log r_e = a \\cdot \\log \\sigma_0 + b \\cdot <\\!\\mu\\!>_e + c. } \\label{FPEQ} \\end{equation} Due to the small intrinsic dispersion ($\\sim 10 \\%$ in \\re \\, and \\s0, and $\\mathrm{\\sim 0.1~mag}$ in \\mie), the FP is considered as a powerful tool to measure galaxy distances and to constrain the processes driving galaxy evolution (e.g. \\citealt{KJM93, JFK96, JOR99, KIV00}).  A relevant projection of the FP is the correlation between \\RE \\, and \\mie, also known as Kormendy relation (hereafter KR): \\begin{equation} \\mathrm{ <\\!\\mu\\!>_e=\\alpha + \\beta \\cdot \\log R_e,   \\label{EQKR} } \\end{equation} where \\RE \\, is the effective radius in $\\mathrm{kpc}$, and $\\beta \\simeq 3$ \\citep{KOR77}. Different studies have demonstrated that in the nearby universe bright ETGs define a sharp sequence in the \\lRE--\\mie \\, plane, with an intrinsic dispersion of $\\sim 0.3$ -- $0.4$ in \\mie \\, (\\citealt{HaK87, HOS87, SaP91, SaL01}). Moreover, \\citet{CCD92} have shown that $\\mathrm{ETGs+bulges}$ form two distinct families in the plane of the effective parameters: that of the bright ETGs, following the KR, and another `ordinary' family, whose properties are more disperse and heterogeneous. Recently, \\citet[see also references therein]{GrG03} demonstrated that the claimed dichotomy between the bright and the dwarf ETGs is only apparent and that there is a continuous structural relation between the two classes. In particular, the different behavior of bright and dwarf ETGs in the \\lRE--\\mie \\, plane, and the change in slope of their relations, do not imply a different formation mechanism but can be interpreted with systematic changes in the profile shape with galaxy magnitude.  In order to understand the processes underlying galaxy formation and evolution, it is crucial to investigate the FP and the other relations between galaxy parameters at different look-back times. In the monolithic formation scenario \\citep{Lar74}, ETGs form at very high redshifts in a strong burst of star formation, and the distribution of galaxies in the space of observed quantities changes with redshift due to the passive fading of their stars.  In the hierarchical framework, bright ETGs form by the progressive merging of smaller units, and the FP relations are built up with redshift.  Different works have shown (1) that (dissipation-less) merging seems to move galaxies along the FP \\citep{CdCC95, DCR02}, and (2) that in hierarchical models, the luminosity--weighted ages of stellar populations in cluster galaxies appear to be the same of those predicted for pure passive evolution (e.g. \\citealt{KAU96}).  All these facts suggest that for both scenarios dynamical effects could play a minor role in the evolution of scaling correlations. It is expected that the FP slopes change with redshift mainly for possible differences in properties of galaxy stellar populations (such as age and metallicity) along the early-type sequence, while the zero point varies consistently with the evolution of a passive stellar population.  In the hierarchical framework, however, it is also expected that the distribution in the space of observed quantities evolves with redshift, as merging proceeds to larger and larger sizes. Clearly, in order to analyze this subject, large samples of galaxies are needed at different look-back times.  In the recent years, it has become feasible to study the scaling relations at intermediate redshifts. Due to the heavy demand of observing time in velocity dispersion measurements, many studies have investigated pure photometric laws, like the KR. The main effort to constrain the evolution of the KR slope with redshift has been performed by \\citet{ZSB99}, who analyzed ETGs properties in four clusters at $\\mathrm{z \\sim 0.4}$ and in one at $\\mathrm{z=0.55}$, with an average of $\\sim 20$--$30$ galaxies per cluster. These authors found that the slope is comprised between 2.2 and 3.5 both for the Coma and for the distant clusters, depending on the selection criteria of the samples, and suggested, therefore, that it could not have evolved since $\\mathrm{z \\sim 0.55}$.  The evolution of the KR zero point for cluster galaxies has been investigated in different works in order to perform the Tolman surface brightness test for the universal expansion \\citep{SaP91, PDdC96, SaL01, LuS01a, LuS01b, LuS01c} and to derive the luminosity evolution of ETGs (e.g. \\citealt{BAS98, ZSB99}). All these studies concluded that the measured change is consistent with what expected on the basis of cosmological dimming plus passive aging of stellar populations formed at high redshifts.  The KR was investigated at high redshifts for galaxies in the Hubble Deep Field by \\citet{FCA98}, who found that ETGs in the field seem to follow the same restframe KR also at redshifts $\\mathrm{z \\sim 2 \\! - \\! 3}$.  As shown by \\citet{LBM02}, it is possible to derive structural parameters of galaxies at intermediate redshifts by using ground--based imaging.  In the present work, we use structural parameters in the V-band restframe for a total of $\\mathrm{N = 228}$ spheroids in the clusters \\a209 \\, at $\\mathrm{z=0.21}$, \\ac118 \\, at $\\mathrm{z=0.31}$, and $\\mathrm{EIS \\, 0048 \\! - \\! 2942}$ (hereafter \\eis0048) \\, at $\\mathrm{z=0.64}$, and investigate the evolution of the distribution in the \\lre -- \\mie \\, plane up to a look-back time of $\\mathrm{\\mathrm{\\sim 6~Gyr}}$. Surface photometry for galaxies in \\a209, \\ac118 \\, and \\, \\eis0048 is published in \\citet{LBM03b}.  One of the main challenges of previous works on the subject has been represented by selection effects, which can critically affect the determination of the KR.  To work out this problem, we introduce a new fitting procedure, which provides un-biased estimates of the KR coefficients, and allows for the first time to investigate the evolution over redshift of the whole distribution in the plane of effective parameters.  The layout of the paper is the following. In Section~\\ref{SDATA}, we present the samples used for the analysis. Section~\\ref{SKREL} describes the new fitting procedure and deals with the evolution of the KR slope and the comparison of the \\lre--\\mie \\, distributions at different redshifts. Section~\\ref{SLEV} shows the luminosity evolution inferred by the KR. Conclusions are drawn in section~\\ref{SCONC}.  In the following, we will assume the cosmology $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$ and $\\mathrm{ H_0= 70~Km s^{-1} Mpc^{-1} }$. With these parameters, the age of the universe is $\\mathrm{\\sim13.5~Gyr}$, and the redshifts $\\mathrm{z=0.21, z=0.31}$, and $\\mathrm{z=0.64}$ correspond to look-back times of $\\mathrm{\\sim 2.5, 3.5}$, and $\\mathrm{6~Gyr}$ respectively. ",
        "conclusions": "\\label{SCONC} We have studied the evolution of the optical, V-band restframe, Kormendy relation in the range $\\mathrm{0 \\la z \\la 0.64}$, by using ground-based effective parameters derived in previous works for the clusters of galaxies \\a209 ($\\mathrm{z=0.21}$), \\ac118 \\, ($\\mathrm{z=0.31}$), and \\eis0048 \\, ($\\mathrm{z=0.64}$), and published data for the Coma cluster ($\\mathrm{z \\sim 0.023}$). Galaxies were selected by the photometric redshift technique for \\ac118 \\, and \\eis0048 \\, and by spectroscopic redshifts complemented with photometric data for \\a209.  We introduce a new fitting procedure, the $\\mathrm{MLS}$ fit, which provides un-biased estimates of the KR coefficients. This allows, for the first time, to investigate the evolution of the slope and of the scatter of the KR, as well as of the whole distribution in the plane of effective parameters.  The results of the present work can be summarized as follows. \\begin{description} \\item[$\\bullet$] Spheroids define a tight sequence in the plane of the effective parameters at least back to $\\mathrm{z \\sim 0.64}$. \\item[$\\bullet$] The slope $\\beta$ and the intrinsic dispersion \\smie \\, turn out to be consistent at all redshifts. The mean values of $\\beta$ and \\smie \\, are $2.92\\pm0.08$ and $0.40\\pm0.03$, respectively. \\item[$\\bullet$] The fact that the KR slopes are the same among all the clusters constrains the possible change in the formation epoch of galaxy stellar populations per decade of radius to be smaller than $\\sim 28 \\%$ for SSP models, and smaller than $\\sim 18 \\%$ for populations having an exponential SFR with $\\mathrm{\\tau=1~Gyr}$. \\item[$\\bullet$] The mean value and the width of the distribution of galaxy sizes are consistent for each pair of clusters. Moreover, the Kolmogorov--Smirnov tests indicate that the whole distribution in the \\lre -- \\mie \\, plane do not vary over redshift. \\item[$\\bullet$] In agreement with previous studies, the evolution of the KR zeropoint is explained by the superposition of the Tolman signal and the passive evolution of stellar populations with high formation redshift, $\\mathrm{z_f>2}$. \\end{description}  These results imply that, whatever the mechanism driving galaxy evolution is, (1) it has already built up the Kormendy relation at $\\mathrm{z=0.64}$ and (2) it does not affect significantly the properties of spheroids in the plane of effective parameters from that redshift down to $\\mathrm{z\\sim0}$.  The fact that the KR slope does not evolve with redshift sets constrains on the differential luminosity evolution along the early-type sequence, and, therefore, on the properties of galaxy stellar populations as a function of galaxy radius.  These results should be taken into account by any model aimed at reproducing the physical processes underlying formation and evolution of spheroids."
    },
    "0212/astro-ph0212011_arXiv.txt": {
        "abstract": "{ We review recent progress by the RSAA team in elucidating the dynamical evolution of the various classes of Narrow Line Regions (NLR) in active galaxies. }   \\listofauthors{M.~A. Dopita, G.~V. Bicknell, R.~S. Sutherland \\& C.~J. Saxton, C.~J.} \\indexauthor{Dopita, M.~A.} \\indexauthor{Bicknell, R.~S.} \\indexauthor{Sutherland, R.~S.} \\indexauthor{Saxton, C.~J.}  \\begin{document} ",
        "introduction": " ",
        "conclusions": "In this paper we have presented observational and theoretical evidence to support the following model for the gross features of the temporal evolution for the NLR:\\\\ 1. Early on, the radio jet (if  there is one), or else a strong outflow of thermal gas from the inner accretion disk pushes a strong shock into the surrounding galactic ISM.\\\\ 2. Dense molecular clouds in the shocked region are crushed by radiative shocks, but along the jet axis, such clouds may be rapidly ablated and shredded by non-radiative shocks. In this phase, the optical/UV spectral signature is one of fast radiative shocks.\\\\ 3. Eventually, the jet escapes from the dense region of the galaxy, and the pressure driving the shocked cocoon drains away. The wall shocks then become first radiative, and finally, momentum conserving. Shock-induced star formation may eventually occur within them, provided they become unstable to their self-gravity. The opening angle of the shocked cocoon provides the ``ionization cone''. This opening angle tends to increase with time, allowing an evolution from a Seyfert II-like appearance early on, to predominantly Seyfert I-like or BLR-like spectra at late times.\\\\ 4. In this phase, the crushed clouds left behind in the cocoon are photoablated by the strong UV field from the central engine, and the ablated ionized gas may be radiatively accelerated in the radial direction. in this phase the NLR presents the signatures of a photoionized plasma, and in Seyfert galaxies the coronal lines may be strong. \\\\"
    },
    "0212/astro-ph0212227_arXiv.txt": {
        "abstract": "Results of calculations of the flow structure in binary systems with mass exchange driven by stellar wind are presented. 2D simulations have been carried out using the Roe-Oscher scheme. The fine grid we used allowed us to detect the details of the flow structure. In particular, it was found out that solutions with the wind velocity $V$ of the order of the orbital velocity of the system $V_{s}$ are unstable: in solutions with $V<V_{s}$ -- the steady accretion disk takes place and if $V>V_{s}$ -- the cone shock forms. Consideration of the minor variations of the wind velocity $V$ around $V_s$ shows that they can influence greatly the flow structure causing transition from disk accretion to the accretion from the flow. During the flow rearrangement period (the accretion disk destruction) the accretion rate increases in dozens times for a short time ($\\sim$ 0.1 of the orbital period of the system). This change can result in the drop-off of the gas from the accretor that is usually associated with activity in these systems. ",
        "introduction": "The overwhelming majority of stars are binaries. In wide pairs the effect of one component to another is negligible but if the components are close enough we can't ignore the interaction between them. One of the most bright manifestations of this interaction is the mass exchange between the components. Such systems are called interacting binaries. The study of the flow structure in these systems is important for the consideration of their evolutionary status as well as for interpretation of observations.  Surely, the strong interaction is observed for close systems where at least one of components fills its critical surface (so called Roche lobe), i.e. equipotential surface passing through the inner Lagrangian point $L_1$ (see, e.g., [\\ref{Kniga}]). In the $L_1$ point the pressure gradient is not balanced by the gravitational force and the component that fills its Roche lobe loses matter through the $L_1$ point.  But in many cases even if no one of components fills its Roche lobes the mass transfer between them may be rather intense due to the stellar wind from one of the components (or in some cases from both of them). In this work we consider the flow structure in a binary system with mass exchange driven by stellar wind using symbiotic stars as the example.  Now it is generally recognized that classic symbiotic stars are binary systems consisting of a cool giant and a hot compact star (white dwarf or subdwarf) surrounded by nebulosity. Symbiotic stars are also characterized by a significant photometric and spectroscopic variability and, in particular, by the irregular flare activity. It is also avowed that components in these systems don't fill their Roche lobes and there are evidences that the giant loses mass at the high rate due to stellar wind and the intense interaction between the components takes place. These hypotheses are also supported by observations (see, e.g., [\\ref{Boyar92}]).  In order to explain the flare activity in symbiotic stars different mechanisms were proposed (see, e.g., [\\ref{MK92}]). In particular, the most probable scenario for a classic symbiotic star is the one where the additional energy release is due to the fluctuations of the accretion rate near the level of the stationary thermonuclear burning [\\ref{FC95}]. If the accretion rate exceeds the maximum value when the hydrogen can burn in a layer source of a degenerate core, the accreting gas stores above the burning layer and grows to the giant's size [\\ref{TU}--\\ref{IbenTu}]. This process is usually associated with outbursts in classical symbiotic stars. What are these changes in the accretion rate caused by? In a number of works (e.g., [\\ref{Boyar92},\\ref{Friedjung93}]) some attempts to explain such a flare activity by the giant's wind variability were made but the observations carried out before the outburst registered only insignificant changes of the wind [\\ref{Friedjung93},\\ref{Kenyon83}]. This fact imposes restrictions to all outburst scenarios caused by changes of giant's characteristics.  In this work we have carried out the 2D numerical investigation of mass transfer in classical symbiotic systems using TVD method that is considered to be the best for numerical simulation of mass transfer in binary systems~[\\ref{Kniga}]. The main attention was paid to consideration of the solution where stellar wind velocity has changed from below to above of the orbital velocity of the system as well as to study of the transition from the disk accretion to the accretion from the flow and appropriate rearrangement of the flow structure. ",
        "conclusions": "\\bigskip  The 2D numerical investigation of the gas flow structure in the interacting binary systems with the mass exchange via stellar wind (and namely for their subclass -- symbiotic systems) has been carried out using the Roe-Oscher scheme. The modelling has been carried out for different values of the wind velocity. The effect of the velocity change on the system's behavior was of special interest.  It is shown that after the wind velocity increase the collision of two codirectional flows results in the formation of two shock waves that begin to propagate from the donor's surface. These shocks can lead to the significant changes in the flow structure near the accreting component and namely to the accretion regime change.  These results were used as the ground for the new mechanism providing the change of the accretion rate and explaining the nature of outbursts in symbiotic stars. The point of this mechanism is in the abrupt change of the flow structure near the accreting component as a result of minor fluctuations of the wind from the mass-losing component. These changes come out in the accretion regime change, namely, in transition from the disk accretion to the accretion from the flow when the wind velocity increases. The process of the flow rearrangement in such cases is followed by the jump in the accretion rate.  Results of modelling of the flow structure in the classical representative of the studied type of binary systems Z~And show that in the quiescent state characterized by the wind velocity from 25 to 40~km/s the steady accretion disk forms in the system. In the given regime the accretion efficiency is $\\sim $1\\% of the donor's mass loss. If the wind velocity increases and becomes greater than 35~km/s the disk disappears and the cone shock wave forms. In stationary solutions for high-velocity regimes ($35-75$ km/s) the accretion rate is slightly higher and equals $\\sim $2\\% of the donor's mass loss. But during the flow rearrangement when the wind with increased velocity reaches disk and crushes it the accretion rate increases in dozens times!  The analysis of the observations of symbiotic stars during the active stages shows that the most part of the registered manifestations can be explained in terms of the drop-off of the optically thick shell by the accretor [\\ref{FC95}]. The suggested mechanism provides the accretion rate jump and corresponding increase of the energy release rate large enough to provide the drop-off of the part of the gas from the accretor. Therefore it leads to the appearance of the flow features that can explain the observations."
    },
    "0212/astro-ph0212361_arXiv.txt": {
        "abstract": "Accurate atomic data have become available in the recent past due to the demands of astrophysics and fusion research. We report on the impact of such data on non-LTE line-formation calculations for CNO in early-type stars. Considerable improvement is achieved by the derivation of consistent results from practically all available spectroscopic indicators, regardless of ionization stage or spin system, and the uncertainties in the analyses are drastically reduced. Moreover, systematic trends are revealed, e.g.~an increase of the N\\,{\\sc i} abundances from previous studies of BA-type supergiants by a factor of two is indicated. The present work promises stringent observational constraints on chemical mixing in the course of massive star evolution. First results on BA-type supergiants in the Galaxy and the Magellanic Clouds are discussed. ",
        "introduction": "Carbon, nitrogen and oxygen represent the body of the heavy elements in the universe. Knowledge of their abundances is a key ingredient for understanding the evolution of stars, galaxies and the universe. However, reliable and accurate information on CNO abundances is scarce. In the late-type stars it is primarily the convective nature of the atmospheres that introduces systematic uncertainty, such that even the solar CNO abundances are under debate still (e.g.~M.~Asplund, these proceedings). On the other hand, non-LTE effects prevail in the earlier spectral types. Besides efficient numerical techniques, accurate atomic data are required for solving the non-LTE radiative transfer problem. Vast amounts of radiative and collisional data from quantum-mechanical {\\em ab-initio} calculations have been provided by the Opacity Project (Seaton et al.~1994) and the IRON Project (Hummer et al.~1993) in the recent years, supplemented by contributions from fusion research and state-of-the-art experiments. Use of the $R$-matrix method in the close-coupling approximation allows for a reduction of the typical uncertainties in the data to $\\sim$10\\% on the mean, in sharp contrast to the order-of-magnitude approximations widely used in astrophysics. In the following we discuss the impact of such data on non-LTE abundance determinations for CNO in early-type stars and report on first applications. ",
        "conclusions": ""
    },
    "0212/astro-ph0212047_arXiv.txt": {
        "abstract": "We have examined the morphological make-up of X-ray bright groups. The brighter galaxies in these groups exhibit clear morphology-density and morphology-radius relations. The group morphology-density relation is offset from the cluster relation in the sense that at a given surface density, X-ray bright groups have a lower spiral fraction. After correcting for projection effects the morphology-3D density relation is still shifted towards fewer spirals for a given 3D density, in comparison with clusters. A simple model which corrects the group data for the effects of projection and for the expected higher merging rate in groups, brings the morphology-density relation into good agreement with that of clusters, suggesting that the relation may be driven by two-body interactions.  The fraction of S0 galaxies in these X-ray bright groups is at least as high as that observed in nearby clusters. Given the low velocity dispersion of groups, this indicates that ram pressure stripping is not the dominant mechanism for S0 formation. ",
        "introduction": "\\label{sec:intro}  The well known relationship between galaxy morphology and density (e.g. \\citealt{dressler80,dressler97}) suggests that the environment of a galaxy may have a significant effect on its evolution. By examining the properties of galaxies as a function of environment it should be possible to gain important insights into the evolution of both galaxies and galaxy systems.  Particularly interesting is the group environment, which is typically made up of between three and a few tens of galaxies. These systems are gravitationally bound systems in which the density and velocities of the member galaxies suggest that mergers and interactions are more common than in clusters or in the field (e.g. \\citealt{mamon92,mamon00}). In addition, the majority of galaxies are found in a group environment \\citep{tully87}, and many galaxies in clusters will once have been part of groups.  Perhaps the most striking connection between group and galaxy properties is the fact that almost all X-ray bright groups contain a bright early-type galaxy located at the spatial and kinematic centre of the group \\citep{zabludoff98}. This central galaxy is presumably the result of early merging activity in the collapsing group core \\citep{governato96} and given the high expected rates of mergers in present day groups, this central galaxy may still be rapidly growing through mergers and accretion. Other galaxies in the group are also likely to be affected by the group environment. \\citet{postman84} and \\citet{ramella99} have shown that galaxies in groups follow a morphology-density relation. However, it is not clear if this is the same relation as for clusters as some X-ray bright groups appear to have spiral fractions consistent with those found in rich clusters \\citep{zabludoff98}.  \\begin{figure*} \\hspace{0cm} \\psfig{file=90percent.eps,angle=-90,width=11cm} \\caption{\\label{fig:spiral_radius}Spiral fraction (including irregulars) versus radius (in units of the virial radius, $R_V$) for X-ray bright groups. Each radial bin contains at least 75 galaxies.} \\end{figure*}  It is also known that there is a connection between the morphological make-up of groups and the presence of diffuse X-ray emission (e.g. \\citealt{pildis95,henry95,mulchaey96}), in the sense that X-ray detected groups tend to be spiral poor, whilst undetected groups contain more late-types. In this and a companion paper \\citep{helsdon02} we aim to look in some detail at the relationship between the X-ray and optical properties of X-ray bright galaxy groups. Here we focus on the morphological makeup of these systems, and examine the morphology-density relation of X-ray bright groups. ",
        "conclusions": "\\label{sec:dis}  Studies of the evolution of the morphology-density relation (e.g. \\citealt{dressler97}) have shown that the fraction of ellipticals in clusters has not evolved significantly since z$\\sim$0.5. The elliptical fractions observed in these groups support the suggestion of \\citeauthor{dressler97} that the formation of ellipticals predates cluster formation, and that they are instead formed primarily in a group environment. However, the S0 fraction in clusters at z$\\sim$0.5 is about 2-3 times smaller than in present day clusters (e.g. \\citealt{dressler97}; although see \\citealt{andreon98} for a different view), suggesting that some process must have transformed the excess spirals seen in these systems into S0s by the present day.  Major mergers do not appear able to produce enough S0s in clusters (e.g. \\citealt{okamoto01}), and a number of other mechanisms have be proposed to explain the transformation of spirals to S0s: ram pressure stripping (e.g. \\citealt{gunn72,quilis00}), galaxy harassment (e.g.\\citealt{moore99}) or unequal mass galaxy mergers (e.g. \\citealt{bekki98}).  The fact that we see a comparable S0 fraction in present day X-ray bright groups and clusters, suggests that whatever process is responsible for converting spirals to S0s in clusters also plays a significant role in creating S0s in these groups.  This would appear to rule out ram pressure stripping of the whole ISM of a galaxy as the dominant process. Indeed, since ram pressure scales as the square of the galaxy velocity within its group or cluster, ram pressure stripping should not be effective in the low velocity dispersion group environment (e.g. \\citealt{zabludoff98,abadi99}). Furthermore, there is also evidence that ram pressure stripping cannot produce sufficient numbers of intermediate bulge-to-disk ratio galaxies in clusters (e.g. \\citealt{okamoto01b,lanzoni00}), where it should be more effective. However, if continuing star formation within spirals is fuelled by continuing infall of gas from a surrounding gas reservoir (still a controversial issue), then such a reservoir should be easily removed by infall into groups as well as clusters, leading to a gradual decline or `strangulation' (\\citealt{larson80,balogh00,balogh00b}) in star formation.  There is observational evidence to suggest that this may be occurring, since late-type galaxies in groups appear to be HI-deficient (e.g. \\citealt{williams87,oosterloo97,verdesmontenegro01}).  The cessation of star formation would not in itself convert spirals to lenticulars, so dynamical interactions are probably involved in any case in effecting structural changes to spirals within dense environments (tidal disruption is needed to suppress spiral features and enhance the relative importance of the bulge).  It is tempting to relate the difference in behaviour in the highest density bin, between the incidence of lenticulars in X-ray bright groups and clusters, to the apparent evolution in cluster S0 content since moderate redshifts. In contrast, the elliptical content of clusters is not found to evolve, and appears to be similar to that of groups after correction for the higher merger rate in groups. It may be, for example, that galaxy harassment within the densest regions of the cluster environment has further boosted the S0 content inherited from precursor groups: galaxy harassment depends on the frequency of encounters \\citep{moore98b}, which will be higher than the merger rate, and should be greatest in the centres of clusters."
    },
    "0212/astro-ph0212271_arXiv.txt": {
        "abstract": "We report observations of the CO \\jseven\\ transition toward the starburst nucleus of NGC~253.  This is the highest-excitation CO measurement in this source to date, and allows an estimate of the molecular gas excitation conditions.  Comparison of the CO line intensities with a large velocity gradient, escape probability model indicates that the bulk of the 2--5$\\rm\\times10^{7}~M_{\\odot}$ of molecular gas in the central 180 pc is highly excited.  A model with T~$\\sim$~120~K, $\\rm n_{H_2}\\sim~4.5\\e{4}~cm^{-3}$ is consistent with the observed CO intensities as well as the rotational \\hh\\ lines observed with ISO.  The inferred mass of warm, dense molecular gas is 10--30 times the atomic gas mass as traced through its [\\ion{C}{2}] and [\\ion{O}{1}] line emission.  This large mass ratio is inconsistent with photodissociation region models where the gas is heated by far-UV starlight.  It is also not likely that the gas is heated by shocks in outflows or cloud-cloud collisions.  We conclude that the best mechanism for heating the gas is cosmic rays, which provide a natural means of uniformly heating the full volume of molecular clouds.  With the tremendous supernova rate in the nucleus of NGC~253, the CR heating rate is at least $\\sim$ 800 times greater than in the Galaxy, more than sufficient to match the cooling observed in the CO lines. ",
        "introduction": "NGC~253 is a nearby ($d \\sim 2.5$~Mpc; \\citealp{mau96} ), nearly edge-on ($i~\\approx~78^{\\circ}$, \\citealt{pen81}), isolated starburst galaxy with the starburst strongly concentrated at the nucleus.  Based on the distribution of the 10--30~$\\mu$m continuum, the bulk of NGC~253's starburst is confined to a $\\sim 60$~pc region centered just south-west of the dynamical nucleus \\citep{tel93}.  The 1--300~$\\mu$m luminosity within about 30\\as\\ is $\\rm 1.6 \\times 10^{10}~L_{\\odot}$ \\citep{tel80} while the total IR luminosity of NGC~253 detected by IRAS is $\\rm L_{\\rm IR} \\approx 2 \\times 10^{10}~L_{\\odot}$ \\citep{ric88}.  There is a clear bar in the 2MASS image of the galaxy, and the kinematics show evidence for orbits in a bar potential (\\citealt{sco85}; \\citealt[and references therein]{das01}).  The bar likely plays a role in transporting gas to the nucleus and thus supporting the starburst.  Current estimates of the neutral gas mass in the central $20'' - 50''$ range from $\\rm 2.5 \\times 10^{7}~M_{\\odot}$ \\citep[hereafter HHR99]{hhr99} to $\\rm 4.8\\times 10^{8}~M_{\\odot}$ \\citep{hou97}.  As is the case in many of the nearby starburst galaxies, much of the molecular gas in NGC~253 appears to be highly excited \\citep{wil92,mao00,war02}.  Observations of \\jfour\\ and \\jsix\\ transitions of CO \\citep{isr02,har91}, as well as \\jone\\ and \\jthree\\ transitions of HCN \\citep{pag95,pag97} suggest that a large fraction of the molecular gas in the nucleus of NGC 253 has $\\rm n_{H_2} > 10^4 cm^{-3}$ and $\\rm T\\sim 100 K$ or higher.  Observations of higher-J CO transitions are critical to measure this warm molecular gas component, and probe its physical conditions.  In our Galaxy, the warm, dense molecular gas conditions which excite the mid-J CO lines (J~=~4--8) are typically associated with molecular cloud surfaces exposed to stellar far-UV radiation.  In these photodissociation regions (PDRs), the UV dissociates and ionizes the gas as it enters the cloud and is absorbed by dust.  The result is a layered geometry based on the extinction of UV; the progression from the exterior inward is ionized, neutral atomic, warm molecular, and finally cool molecular gas on the interior \\citep{th85,wol90,kau99}.  Several of these regions have been studied in star formation sites within the Galaxy \\citep{har85,sch89,sta93}, and they have been discussed in the context of starburst nuclei \\citep{mao00,car94}.  Other potential molecular gas heating sources are outflows and cloud-cloud collisions in which shocks convert dynamical energy to thermal energy of the gas, such as in the circumnuclear disk of the Galactic Center, and in the massive star formation sites W~51 and DR~21 \\citep{jhg87,jaf89}.  Finally, it has been proposed that much of the molecular gas heating in starburst nuclei may be due to cosmic rays \\citep[hereafter SAH]{suc93}. To better understand the energy balance of the molecular gas in the starburst nucleus of NGC 253, we have observed the CO \\jseven\\ transition with our submillimeter imaging spectrometer, SPIFI on the JCMT. ",
        "conclusions": "We have detected bright CO \\jseven\\ emission from the nucleus of NGC~253.  We combine our observations with lower-J measurements of both $\\rm ^{12}CO$ and $\\rm ^{13}CO$ and an LVG model to estimate the CO column density and physical conditions of the molecular gas in the starburst nucleus.  The column density of CO we infer based on the transitions above \\jone\\ is slightly larger than but similar to the estimates made by HHR99 for the central 23\\as.  However, we find that the excitation conditions of this gas must be very high to reproduce the observed \\jsix\\ and \\jseven\\ intensities.  A single component with $\\rm T \\sim 120\\, K$, and $\\rm n\\sim 4.5\\e{4}\\, cm^{-3}$ provides the best fit to all the observed transitions above \\jone\\ for both $^{12}$CO and $^{13}$CO, and is consistent with the ISO molecular \\hh\\ observations.  We adopt this as the favored model.  While a smaller additional mass of cool gas is required to account for the $^{13}$CO \\jone\\ intensity, and more cool gas may be present, the CO column density of warm molecular gas cannot be much less than our derived values.  The large mass of warm gas is very difficult to explain with a photodissociation region scenario.  The molecular gas mass is a factor of 10--30 greater than the atomic gas mass, while the two should be comparable if heated with stellar UV as in the Galactic PDR templates.  Dynamical heating of molecular gas from shocks is also unlikely based on energetic considerations, and given the lack of evidence for significant heating in fast shocks.  On the other hand, the very large cosmic ray flux in the starburst nucleus of NGC~253 readily provides the necessary energy to heat the gas, and a means to deliver energy throughout the volume of large, UV-shielded clouds.  It is likely the dominant heating term in the molecular ISM energy budget.  While some of the warm molecular gas that we observe is likely heated by the far-UV radiation and shocks, our results strongly support the conjecture by \\citep{suc93} that the bulk of molecular gas in starburst nuclei is held at high excitation by the enhanced cosmic ray flux."
    },
    "0212/astro-ph0212337_arXiv.txt": {
        "abstract": "The goal of finding and characterizing habitable planets in other solar systems represents one of humanity's greatest scientific challenges. NASA and ESA have initiated studies of missions that could accomplish this goal within the next ten years. What precursor knowledge do we need before we can initiate such a mission? How large should the first steps be in a program whose ultimate aim is to detect life on other planets? This talk describes different concepts for NASA's Terrestrial Planet Finder and discusses potential precursors in a program that balances scientific return, technological advance, and programmatic risk. ",
        "introduction": "The public and the scientific response to NASA's search for habitable planets and life has been very enthusiastic.  Most recently, the National Academy of Science's Decadal review  of Astronomy and Astrophysics (2000) endorsed the Origins goals,  noting:  \\begin{quote} ``Key problems that are particularly ripe for advances in the coming decade\\dots [include] studying the formation of stars and their planetary systems, and the birth and evolution of giant and terrestrial planets.\" \\end{quote}  \\begin{quote} ``Search for Life outside of earth and, if it is found, determine its nature and its distribution in the galaxy\\dots [This] is so challenging and of such importance that it could occupy astronomers for the foreseeable future.\" \\end{quote}  But the NAS panel added a caveat to its recommendation that NASA proceed with the Terrestrial Planet Finder mission (TPF):  \\begin{quote} ``The committee's recommendation of this mission is predicated on the assumptions that TPF will revolutionize major areas of both planetary and non-planetary science, and that, prior to the start of TPF, ground- and space-based searches will confirm the expectation that terrestrial planets are common around solar type stars.\" \\end{quote}   This talk addresses what we must learn about the prevalence (or absence) of Earth-like planets before starting TPF and discusses what missions might play a role in a balanced program to advance  our scientific knowledge and technological prowess to the level  needed for TPF. ",
        "conclusions": ""
    },
    "0212/nucl-th0212059_arXiv.txt": {
        "abstract": "In interpreting the SNO experiments, accurate estimates of the $\\nu d$ reaction cross sections are of great importance. In our recent work \\cite{NETAL}, we have improved our previous calculation by updating some of its inputs and by incorporating the results of a recent effective-field-theoretical calculation. The new cross sections are slightly ($\\sim$1 \\%) larger than the previously reported values. It is reasonable to assign 1\\% uncertainty to the $\\nu d$ cross sections reported here; this error estimate does {\\it not} include radiative corrections. ",
        "introduction": "The establishment of the Sudbury Neutrino Observatory (SNO) has greatly increased the necessity of detailed theoretical studies of the neutrino-deuteron ($\\nu d$) reactions\\cite{NETAL,NSGK,ando,kn,BCK}. At SNO, neutrino oscillations can be directly tested by measuring the solar electron-neutrino flux using the charged-current (CC) reaction ($\\nu_e d\\rightarrow e^-pp$) and the total solar neutrino flux using the neutral-current (NC) reaction, $\\nu_x d\\rightarrow \\nu_xpn$ ($x=e,\\mu\\ {\\rm or}\\ \\tau$). The recent SNO results\\cite{ahmad} have given definitive evidence for neutrino oscillations. To make the best use of the existing and future SNO data, it is highly desirable to further improve the theoretical estimates of the $\\nu d$ cross sections.  A traditional approach for describing nuclear electroweak processes consists in evaluating the contributions of 1-body impulse-approximation (IA) operators and 2-body exchange-current (EXC) operators defined in the Hilbert space of non-relativistic nuclear wave functions, with the EXC terms derived from one-boson exchange diagrams. We refer to this approach as SNPA (standard nuclear physics approach). The calculations of the $\\nu d$ cross sections (\\signud)\\ based on SNPA have been done by several authors \\cite{kn}; the most recent work due to Nakamura {\\it et al.}\\ shall be referred to as NSGK \\cite{NSGK} and NETAL \\cite{NETAL}.  An alternative approach, effective field theory (EFT), has been applied to the $\\nu d$ reaction by Butler {\\it et al.}(BCK)\\cite{BCK}. The EFT lagrangian used by BCK involves one unknown low-energy constant (LEC), $L_{1A}$, which BCK adjusted to optimize fit to the \\signud\\ of NSGK. After this fine-tuning, the results of BCK were found to agree with those of NSGK within 1\\% over the entire solar-neutrino energy region. Very recently, Ando {\\it et al.} have performed an EFT-motivated calculation of \\signud\\ \\cite{ando} using a method called EFT*; in this approach, originally proposed by Park et al.\\cite{PKMR}, the electroweak transition operators are derived with a cut-off scheme EFT, while the initial and final nuclear wave functions are obtained with the use of a realistic phenomenological $NN$ potential. The EFT* lagrangian contains an unknown LEC, denoted by $\\hat{d}^R$, which plays a role similar to $L_{1A}$ in BCK. In EFT*, however, $\\hat{d}^R$ can be determined directly from the tritium $\\beta$-decay rate \\G3H, which allows a parameter-free calculation of the $\\nu d$ cross sections.  We give here a concise account of the latest SNPA calculation of the $\\nu d$ reaction carried out in NETAL \\cite{NETAL}. The main points of improvement over NSGK are as follows. First, the updated value of the axial coupling constant $g_A$ is used. Secondly, the treatment of the axial-vector exchange current (\\AEXC) is improved. Since the $\\nu d$ reaction in the solar-neutrino energy region is dominated by the Gamow-Teller (GT) transition, \\AEXC\\  is a crucial ingredient that controls the accuracy of calculated cross sections. Among the various terms in \\AEXC, the $\\Delta$-excitation current is the most important one. NETAL uses the $\\Delta$-excitation current determined in Ref. \\cite{Schiavilla} to reproduce the experimental value of \\G3H. Thirdly, NETAL employs the effective Fermi coupling constant $G_F^{\\, \\, \\prime}$ that includes inner-radiative corrections instead of $G_F$ used in NSGK. Furthermore, the stability of theoretical estimates is investigated by comparing the new SNPA calculation with Ando {\\it et al.}'s EFT* calculation \\cite{ando}. ",
        "conclusions": "The calculated total cross section for the CC reaction (\\signudCC)\\ is shown in Fig.\\ref{fig_tot} as a function of the incident neutrino energy ($E_{\\nu}$). Although the cross section is dominated by the GT transition leading to the final ${}^1S_0$ $NN$ state, the transitions leading to higher partial waves are not totally negligible; they account for 1\\% and 4\\% of the reaction rate at $E_\\nu \\sim 10$ MeV and $E_\\nu \\sim 20$ MeV, respectively.  \\begin{figure}[h] \\begin{center} \\includegraphics[width=210pt]{tot.eps}\\\\[-6mm] \\caption{Total cross sections for the CC reaction. The solid- (dotted-) curve includes all (${}^1S_0$) partial waves of the final two-nucleon system.} \\label{fig_tot} \\end{center} \\end{figure}   \\footnotesize \\begin{table}[h] \\begin{minipage}[t]{83mm} \\begin{tabular}[t]{cccc}\\hline $E_{\\nu}$ (MeV) & NETAL & NSGK & normalized \\\\\\hline 5& 1.019& 1.050& 1.020\\\\ 10& 1.023& 1.059& 1.024\\\\ 20& 1.029& 1.068& 1.030\\\\\\hline &&&\\\\[-2mm] \\end{tabular} \\caption{Contributions of EXC for the CC reaction. The ratio, \\signudCC(IA+EXC)/\\signudCC(IA), is shown for NETAL (2nd column) and for NSGK (3rd column). The fourth column corresponds to the use of the ``normalized'' \\AEXCN\\  explained in the text .} \\label{tab_exc} \\end{minipage} \\hspace{7mm} \\begin{minipage}[t]{70mm} \\begin{tabular}[t]{ccc}\\hline $E_{\\nu}$ (MeV) & $\\nu_e d\\rightarrow e^-pp$ & $\\nu d\\rightarrow \\nu pn$\\\\ \\hline 5& 1.003& 1.004\\\\ 10& 1.001& 1.003\\\\ 20& 0.998& 1.001\\\\\\hline &&\\\\[-2mm] \\end{tabular} \\caption{Comparison of SNPA and EFT*. The ratio \\signud(EFT*)/\\signud(SNPA) is shown for the CC and NC reactions.} \\label{tab_eft} \\end{minipage} \\end{table} \\normalsize The ratio \\signudCC(IA+EXC)/\\signudCC(IA) is shown in Table \\ref{tab_exc} for both NETAL and NSGK. In NETAL, the EXC contribution to \\signudCC\\ is $\\sim$2\\%, which is $\\sim$3\\% smaller than the NSGK results. The strength of \\AEXCN\\  was determined by analyzing the $np \\rightarrow \\gamma d$ reaction and assuming a quark-model relation between the axial vector and vector coupling constants for the $N\\Delta$ transition. We note, however, that \\G3H should give a much more direct constraint on \\AEXC\\  than $\\sigma(np\\rightarrow d\\gamma)$, the latter being governed by the vector current. One way to gauge the sensitivity of \\signud\\  to different choices of \\AEXC\\ is to normalize the strength of \\AEXCN\\ near threshold to reproduce \\AEXCNetal. With this normalization applied, the difference between the two EXC models is reduced to 0.2\\%; see the fourth column in Table 1. Thus it is the overall strength of \\AEXC\\ that controls the low-energy $\\nu d$ reactions; \\signud\\ is insensitive to the detailed structure of \\AEXC\\  once EXC is adjusted to reproduce \\G3H.   Comparison with an EFT* calculation \\cite{ando} provides a further test of the reliability of \\signud\\ obtained in NETAL. It is sufficient to make this comparison for the reaction rate leading to the final $^1S_0$ state. Table \\ref{tab_eft} shows the ratio  \\signud(EFT*)/\\signud(SNPA), from which we can conclude that SNPA and EFT* give identical results within 1\\% accuracy. This agreement proves the robustness of the theoretical estimates of \\signud\\  obtained in NETAL.  With the improved inputs as described above, NETAL have obtained \\signud's that are slightly larger ($\\sim$1\\%) than those of NSGK. The \\signud's in NETAL are considered to be reliable within 1\\% accuracy.  Besides the absolute value of \\signud, the ratio $R=$\\signudNC$/$\\signudCC\\ is an important quantity for SNO experiments. Comparing the $R$'s obtained in NETAL, NSGK and EFT*, we can assign $0.5$\\% accuracy to the calculated $R$.  The radiative corrections are expected to affect  \\signud\\ at the level of a few per cent. NETAL only took into account a part of the radiative correction incorporated into $G_F^{\\, \\, \\prime}$; the most recent estimation of the remaining radiative corrections can be found in Ref.\\ \\cite{kurylov}.  A new experimental value of $g_A$ has been presented at this conference by Abele \\cite{newgA}. The reported value, $g_A$=1.274, is significantly larger than the current PDG value used by NETAL. We remark, however, that if the calculation of NETAL is repeated with the use of this new value of $g_A$, the resulting \\signud's would be essentially unchanged. Since the strength of \\AEXC\\  in NETAL is determined to reproduce \\G3H, an increase in $g_A$ is largely compensated by a decrease in \\AEXC, leaving \\signud's essentially unaffected.  This work is supported in part by the Japan Society for the Promotion of Science, Grant No. (c) 12640273, and by the US National Science Foundation, Grant No. PHY-9900756 and No. INT-9730847."
    },
    "0212/astro-ph0212451_arXiv.txt": {
        "abstract": "The NUclei of GAlaxies (NUGA) project is a combined effort to carry out a high-resolution ($<$1$\\arcsec$) interferometer CO survey of a sample of 12 nearby AGN spiral hosts, using the IRAM array.  We map the distribution and dynamics of molecular gas in the inner 1\\,kpc of the nuclei with resolutions of $\\sim 10-50\\,{\\rm pc}$, and study the mechanisms for gas fueling of the different low-luminosity AGN. First results show evidence for the occurrence of strong $m=1$ gas instabilities in Seyferts. NUGA maps allow us to address the origin/nature of $m=1$ modes and their link with $m=2$ modes and acoustic instabilities, present in other targets. ",
        "introduction": "The study of interstellar gas in the nuclei of galaxies is key for understanding nuclear activity and circumnuclear star formation. Within the central kiloparsec, most of the gas is in the molecular phase, which makes CO lines the optimal tracers of nuclear gas dynamics. Up to now, however, CO surveys of galaxies made with single-dish telescopes or interferometers were hampered by insufficient spatial resolution. Furthermore, until very recently (Baker 2000; Jogee et al 2001), the published survey samples have included very few AGN.  NUGA aims at filling this gap by carrying out a high resolution/sensitivity CO survey of a moderately large sample of 12 nearby AGN which span the whole sequence of activity types (Seyf 1, 2 and LINERs). We seek to determine the molecular gas distributions and kinematics in the central 1\\,kpc of the nuclei with resolutions $<$1$\\arcsec$ ($<$10-50$\\,{\\rm pc}$), and to study the different mechanisms driving gas infall to the AGN. Observations in the 2--1 and 1--0 lines of $^{12}$CO are carried out with the IRAM interferometer. On the modeling side, N-body simulations are performed to analyse the evolution of stellar/gaseous gravitational instabilities in the different study cases. NUGA has a multiwavelength approach: the sample is defined based on the availability of high-quality optical and near-infrared images of the galaxies, both from ground-based telescopes and the {\\it HST}; these are essential to study how star formation proceeds in the nuclear region. The long term aim is to complete a supersample of 25-30 objects observed by consortium members, within and outside NUGA, with the IRAM array.  \\begin{figure}  \\plotfiddle{gburillo_fig1.eps}{15.78cm}{0}{65}{65}{-185.7}{-25.2} \\caption{$^{12}$CO(1--0) (left) and $^{12}$CO(2--1) (right) NUGA images, discussed in text, for NGC\\,4826 (Garc\\'{\\i}a-Burillo et al. 2002), NGC\\,7217 (Combes et al. 2002), NGC\\,1961 (Baker et al. 2002), and  NGC\\,3718 (Krips et al. 2002). CO maps unveil a large variety of gravitational perturbations: m=1 modes, m=2 modes and stochastic patterns.}  \\end{figure} ",
        "conclusions": "Preliminary results based on still incomplete observations of 8 galaxies of the sample reveal a wide range of gravitational instabilities in the central 1\\,kpc of NUGA targets. Most remarkably, $m=1$ modes appear at different spatial scales in some AGNs, although these modes might not {\\it universally} favor AGN feeding. Point-symmetric $m=2$ perturbations and stochastic patterns dominate the response in others. Although general conclusions will have to wait for a global analysis of the 30-galaxy supersample, these results already indicate that the correlation between activity type (Seyf 1, 2 or LINERs) and nuclear morphology of the host might be weak at scales of ten-to-hundred pc. If confirmed, the lack of a clear evolutionary scenario may reflect a mismatch between the episodic AGN duty cycle, and the larger time-scales needed to build up nuclear gravitational instabilities which are directly or indirectly involved (by the onset of nuclear starbursts) in fueling AGN."
    },
    "0212/astro-ph0212498_arXiv.txt": {
        "abstract": "First results of X-ray source population studies in \\linebreak[4] nearby galaxies show the potential of \\xmm\\ observations. I will report on first \\xmm\\ \\andro\\ results and on three of our \\xmm\\ projects, an X-ray source population study in the Magellanic Clouds (MCs), a deep raster survey of \\m33, and an investigation of the hot interstellar medium (ISM) in the halo of edge-on galaxies. \\xmm\\ results on several other galaxies and sources within are presented by other authors in these proceedings.  Our MC study is build up of deep pointings probing MC sources down to $10^{33}$~erg~s$^{-1}$\\ and shallower pointings to confirm candidates from our ROSAT derived lists of X-ray binaries (XRBs), super-soft sources (SSSs), and supernova remnants (SNRs). First \\xmm\\ detections of a 455 s pulsar in the Small Magellanic Cloud (\\smc) and the results of the Large Magellanic Cloud (\\lmc) deep field confirm the validity of our strategy.  Our \\m33\\ raster pointing aims for luminosities as low as $10^{35}$~erg~s$^{-1}$, a factor of 10 below the sensitivity limit of the ROSAT observations. The survey will allow us to characterize the sources using extent, spectra, hardness ratios and time variability to build  up a  unprecedented census of the X-ray source content of \\m33. Of specific interest are the active source in the nuclear area and the diffuse emission in the inner disk.  \\xmm\\ observations of the active galaxy \\linebreak[4] \\ngc3079 and of the starburst galaxy  \\n253 are used to characterize the point-like sources and the hot ISM in the disk and from the halo of these galaxies. ",
        "introduction": "Observations of the \\ein\\ and ROSAT observatories demonstrated that the X-ray emission of nearby late type galaxies (spirals and irregulars) can be separated in contributions from point-like sources (low- and high-mass XRBs, SSS, young supernovae, SNRs, etc.) and diffuse emission (hot ISM within the disk, from an outflow, or from the galaxy halo). In some of the galaxies X-ray emission from an active nucleus (AGN) could be detected (sometimes heavily absorbed) dominating the galaxy's X-ray emission. The imaging proportional counters aboard \\ein\\ (IPC) and ROSAT (PSPC) were able to spatially resolve these components. However, the energy resolution was only sparse. The CCD detectors aboard the ASCA and BeppoSAX observatories on the other hand provided much better energy resolution and a broader energy coverage. However with the poorer spatial resolution, in most cases it was not possible to resolve the different emission components. The average parameters extracted for the galaxies as a whole, are of rather limited value.  With the new generation X-ray observatories \\chandra\\ and \\xmm, we now have the sensitivity and spatial and spectral resolution to probe significantly deeper. The high throughput of \\xmm\\ combined with the  spatial resolution comparable to the ROSAT high resolution imager (HRI) allows us in X-ray source population studies (1) to detect sources more than a factor of 10 deeper than before; (2) to classify the sources in a broad X-ray energy band (spectra, hardness ratios, variability, extent, pulsations); (3) to determine  log N / log S distributions. For the hot ISM, we are able (1) to resolve it from X-ray point sources; (2) to analyze the spatial variability of its X-ray spectrum with medium and high energy resolution. With the results for different galaxies, we can search for dependencies of the X-ray parameters on galaxy type, metallicity, and star forming history.  In the following I will report on \\xmm\\ results that demonstrate these possibilities. Interesting additional new \\xmm\\ results on \\andro, \\ng300\\ and other nearby galaxies (and individual sources within) will be presented in these proceedings. ",
        "conclusions": "The first observations of nearby late type galaxies have demonstrated the superb performance of the XMM-New\\-ton instruments. However, sometimes high instrument \\linebreak[4] background causes problems by strongly reducing the useful observation time below the times derived in the feasibility calculations and scheduled for observation. Another problem for homogeneous galaxy observations is imposed by the very stable \\xmm\\ attitude control system which does not smear out small scale detector structures  like CCD gaps or bad pixels. One therefore might argue for the introduction of a satellite dithering mode.  Besides these small obstacles, the big advantages of \\xmm\\ for nearby galaxy studies compared to earlier observatories are the \\begin{itemize} \\item good spatial resolution and high collection power for EPIC and RGS spectra \\item high detection potential for XRBs, X-ray pulsars, SNRs, CVs \\item detectability of ISM and especially soft halo emission ($kT \\approx 0.15$ keV) in low background EPIC PN thin filter observations \\item advantage of simultaneous multi-instrument data collection \\end{itemize}  The first \\xmm\\ observations promise a wealth of exciting results when new \\xmm\\ nearby galaxy observations are performed and analyzed. The interpretation will certainly benefit when high spatial resolution \\chandra\\ observations of the galaxies are combined with the collecting power of \\xmm."
    },
    "0212/astro-ph0212417.txt": {
        "abstract": "Cosmic shear measurements have now improved to the point where they deserve to be treated on par with CMB and galaxy clustering data for cosmological parameter analysis, using the full measured aperture mass variance curve rather than a mere phenomenological parametrization thereof. We perform a detailed 9-parameter analysis of recent lensing (RCS), CMB (up to Archeops) and galaxy clustering (2dF) data, both separately and jointly. CMB and 2dF data are consistent with a simple flat adiabatic scale-invariant model with $\\Ol=0.72\\pm 0.09$, $h^2\\Oc=0.115\\pm 0.013$, $h^2\\Ob=0.024\\pm 0.003$, and a hint of reionization around $z\\sim 8$. Lensing helps further tighten these constraints, but reveals tension regarding the power spectrum normalization: including the RCS survey results raises $\\sigma_8$ significantly and forces other parameters to uncomfortable values. % raising it from $\\sigma_8=0.74\\pm 0.06$ to $\\sigma_8=0.92\\pm 0.04$ %by forcing other parameters to dubious values. Indeed, $\\sigma_8$ is emerging as the currently most controversial cosmological parameter, and we discuss possible resolutions of this $\\sigma_8$ problem. We also comment on the disturbing fact that many recent analyses (including this one) obtain error bars smaller than the Fisher matrix bound. We produce a CMB power spectrum combining all existing experiments, and using it for a ``MAP versus world'' comparison next month will provide a powerful test of how realistic the error estimates have been in the cosmology community. % RELATED TO THE ANOMALOUSLY LOW chi2? ",
        "introduction": "%% mention: 1. best CMB+LSS analysis so far - but error bars too small, %          2. worthwhile to treat lensing, lensing does help %          3. there is tension, norm too high for CMB+LSS  An avalanche of precision data has revolutionized our ability to constrain cosmological models and their free parameters in recent years \\cite{Lange00,boompa,Balbi00,observables,Jaffe00,Padmanabhan00,Lineweaver00,10par,concordance,Kinney01,Hannestad01,Hannestad01b,Griffiths01,Phillips01,BondConfProc,CosmicTriangle,Novosyadlyj00,Novosyadlyj00b,Durrer00,Bridle00,Turner01,Holder01,Efstathiou01,Pryke01,Stompor01,qmap3,KnoxPage00,VSA02d,CBI02a,xiaomin01,Tegmark02,Lewis02,Venkatesan02,Melch02a,Melch02b,Melch02c}. % ANY MORE PARAMETER PAPERS WORTHH CITING ON THE FLOLLOWING LISTL? %\\cite{Netterfield01,Halverson01,Lee01,Padin01,qmask,qmaskpow,Netterfield97,qmap1,qmap2,Bennett96,CBI02,CBI02b,VSA02a,VSA02b,VSA02c,Bunn96a,Bunn96b} Around 1998, cosmic microwave background (CMB) data had become good enough to allow a realistic high-dimensional parameter space to be explored by fitting theoretical models directly to the measured CMB power spectrum (\\eg, \\cite{10par,Lineweaver00}). % \\rg\\nnn Lineweaver C H;1998;ApJL;505;L69;L98 % {astro-ph/9805326 (``L98'')} % Title: The Cosmic Microwave Background and Observational Convergence in the Omega_m - Omega_lambda Plane % Author: Charles H. Lineweaver (UNSW,Sydney) % Comments: This version conforms to the version accepted by the Astrophysical Journal Letters. Minor revisions include references, % typos and phrasing. 13 pages including 3 Figures % The Astrophysical Journal, Volume 505, Issue 2, pp. L69-L73. In contrast, other cosmological constraints were included merely as priors on individual parameters, for instance the baryon density from Big Bang nucleosynthesis, the matter density from galaxy cluster abundance and the Hubble parameter from direct observations.  With the advent of improved galaxy redshift surveys like PSCz \\cite{Saunders00}, 2dF \\cite{Percival01} and SDSS \\cite{York00}, %\\rf\\nn York D {\\etal};2000;AJ;120;1579 % astro-ph/0006396 %     Title: The Sloan Digital Sky Survey: Technical Summary %     Authors: D.G. York, et al. (The SDSS Collaboration) %     Comments: 9 pages, 7 figures, AAS Latex. To appear in AJ, Sept 2000 the galaxy power spectrum $P(k)$ was upgraded and admitted to the ``CMB club'': it was so accurately measured that people started fitting models directly to the measured $P(k)$ curve rather than merely to some phenomenological parametrization thereof (say in terms of an amplitude and a ``shape parameter'') \\cite{10par,concordance,Efstathiou01}. The Lyman $\\alpha$ forest (Ly$\\alpha$F) power spectrum has now undergone the same upgrade \\cite{Hannestad01,xiaomin01}  One promising cosmological probe is still conspicuously absent from this club and has so far been left out in the cold: the cosmic shear power spectrum measured with weak gravitational lensing. Like the CMB, it has the advantage of probing the dark matter distribution directly \\cite{Hoekstra02b}, without murky bias issues. Despite the spectacular progress since this elusive signal was first detected in 2000 \\cite{Waerbeke00,Bacon00,Wittman00,Kaiser00}, increasing the sky area covered by almost two orders of magnitude \\cite{Hoekstra02b,RCS01a,RCS01b,Bacon02,Waerbeke01,Hoekstra02a,Refregier02,Waerbeke02,Jarvis02}, it has so far only been included in cosmological parameter studies as a prior on two parameters: the power spectrum normalization and the matter density. Yet it has been shown that the detailed shape of the weak lensing power spectrum $P_\\kappa(\\l)$ potentially contains almost as much cosmological information as the CMB \\cite{HuOkamoto02,Hu98,Hu00}, with great prospects for degeneracy breaking and cross-checking.   %With the newly in-date experimental data,\\eg,  CBI and VSA for %CMB,\\cite{Padin01,CBI02,CBI02b,VSA02a,VSA02b,VSA02c} %2dF\\cite{Lahav01,Colless01,Percival01,Tegmark2df01} for galaxy clustering and %RCS\\cite{RCS01a,RCS01b,Hoekstra02a} for weak lensing,  it is worthwhile for us to   The goal of this paper is to treat weak lensing on an equal footing with CMB and galaxy clustering, fitting theoretical models directly to the CMB, LSS and lensing power spectra. Our results can also be viewed as the last stand before MAP: using our constraints combining the state-of-the-art available data for a ``MAP versus world'' comparison next month will provide a powerful test of how realistic the error estimates and underlying assumptions have been in the cosmology community.  The rest of this paper is organized as follows. We present our analysis method in \\sec{MethodsSec}, describe the CMB, galaxy and lensing data used in \\sec{DataSec}, compute quantitative model constraints in \\sec{ParameterSec} and discuss our conclusions in \\sec{DiscussionSec}. ",
        "conclusions": "\\label{DiscussionSec}  Since cosmic shear measurements have now improved to the point where they deserve to be treated on par with CMB and galaxy clustering data, we have perform a detailed 9-parameter analysis of recent lensing, CMB and galaxy clustering data.  For the CMB part, we constructed and used a data set combining all available data taking calibration errors and beam uncertainties into account, effectively calibrating them off of each other. Using this combined power spectrum for a ``MAP versus world'' comparison next month will provide a powerful test of how realistic the error estimates have been in the CMB community.  Even though 2002 has added substantial amounts of information on the CMB power spectrum, \\eg, from CBI and Archeops, the combined CMB and 2dF data remains comfortably consistent with a simple flat adiabatic scale-invariant concordance model with $\\Ol=0.72\\pm 0.09$, $h^2\\Oc=0.115\\pm 0.013$ and $h^2\\Ob=0.024\\pm 0.003$. The new data has added a hint of reionization around $z\\sim 8$. These constraints assumed negligible contributions from spatial curvature and massive neutrinos.  Including the RCS lensing data in our analysis confirmed with real data what was predicted by Fisher matrix forecasts \\cite{Hu98}: that lensing has the statistical power to substantially help with cosmological parameter constraints in two ways: by breaking degeneracies and by providing cross-checks. Intriguingly, one of the cross-checks failed, so let us now turn to this failure in more detail.  \\subsection{The $\\sigma_8$ problem}  Our results suggest that the tension between the lensing and other data comes from the overall normalization of the power spectrum. We saw that the CMB+2dF data preferred $\\sigma_8=0.74$, in good agreement with results from last year \\cite{Lahav01}. In contrast, the RCS lensing data prefers higher values \\cite{Hoekstra02a} and pushed the normalization up to $\\sigma_8=0.92$ in our joint analysis.  % PHYSICAL INTUITION AS TO WHAT'S GOING ON: Physically, it is easy to see why this pushed other parameters in the way it did, notably increasing $\\tau$ and decreasing $\\Ol$. If we increase the value of $\\sigma_8$, the CMB and galaxy power spectra will both increase in amplitude. This causes no problems for fitting the 2dF data, since our method can merely readjust their bias, but causes serious difficulties with the CMB. %Hoekstra et al.\\cite{Hoekstra02a} pointed out that Thomson scattering of CMB photons lead to a degeneracy between %$\\sigma_8$ and $\\tau$. Yes. But many others pointed it out earlies, so it's safest to cite nobody here. The acoustic peaks can be brought back down to acceptable heights by increasing $\\tau$, but this provides only a partial solution to the problem, since it leaves the power on very large $\\l\\simlt 10$ scales larger than observed by COBE DMR. Lowering $\\Ol$ helps with this, since it increases the linear growth factor and thereby lowers the entire CMB power spectrum for fixed current matter clustering. $\\Ol$ can only be lowered by a moderate amount, however, since the matter density $\\Om\\sim 1-\\Ol$ will increase and give the galaxy power spectrum the wrong shape. The response of the parameter fitting algorithm is therefore to do a little bit of both tweaks, increasing $\\tau$ and lowering $\\Ol$. Indeed, similar results were obtained in a smalled parameter space by Hoekstra {\\etal} \\cite{Hoekstra02a}, who found that they could only reconcile their lensing normalization with CMB for a rather large reionization optical depth $\\tau=0.12\\pm0.04$.  To verify that this is the correct physical interpretation of what is going on, we repeated our entire analysis with the amplitude of the RCS data lowered by a factor of $(0.7/0.86)^2$, which is roughly the ratio of normalizations preferred by CMB+2dF and RCS respectively. This analysis gave results fully consistent with those from CMB+2dF alone, for example, $\\tau=0.05\\pm0.03$, $\\Ol=0.76\\pm0.13$, $\\od=0.11\\pm0.02$, $\\ob=0.024\\pm0.005$, $\\ns=1.00\\pm0.04$, $\\As=0.52\\pm0.10$, $r=0.10+0.23$, $b=0.82\\pm0.09$, $z_{ion}=7.64\\pm3.18$ and $\\sigma_8=0.70\\pm0.11$. These results not only demonstrate consistency, but also show that by combining lensing and CMB and LSS, we can get stronger constraints on parameters like $\\Ol$, $\\ns$, $\\sigma_8$ and $z_{ion}$, just as had been predicted using information theory \\cite{Hu98}.  \\def\\sss{\\hglue5mm} \\bigskip \\noindent %\\fcolorbox{blue}{yellow} {\\footnotesize {\\bf Table 3} -- Recent measurements of the power spectrum normalization $\\sigma_8$. For results quoted as fitting formulas involving $\\Om$, we have set $\\Om=0.3$. Error bars are $1\\sigma$ --- for \\cite{Hoekstra02a} and \\cite{Jarvis02}, we have simply divided their quoted 95\\% errors by 2. \\bigskip \\begin{center} {\\footnotesize \\begin{tabular}{lll} \\hline Analysis&&$\\sigma_8$\\\\ \\hline {\\bf Clusters:}\\\\ %Eke, Cole \\& Frenk (1996)      &$0.93\\pm 0.07$\\\\ %Carlberg {\\etal} (1997)        &$0.82\\pm 0.03\\$\\ %Fan \\& Bahcall (1998)          &$1.18^{+0.24}_{-0.22}$\\\\ %Pen {\\etal} (1998)             &$1.00\\pm 0.09$ \\sss Pierpaoli {\\etal} (2001)        &\\cite{Pierpaoli01}     &$1.02^{+0.07}_{-0.08}$\\\\ \\sss Borgani {\\etal} (2001)          &\\cite{Borgani01}       &$0.76^{+0.08}_{-0.05}$\\\\ \\sss Reiprich \\& B\\\"ohringer (2001)  &\\cite{Reiprich01}      &$0.68^{+0.08}_{-0.06}$\\\\ \\sss Seljak {\\etal} (2001)           &\\cite{Seljak01}        &$0.75\\pm 0.06$\\\\ \\sss Viana {\\etal} (2001)            &\\cite{Viana01}         &$0.61\\pm 0.05$\\\\ \\sss Bahcall {\\etal} (2002)          &\\cite{Bahcall02}       &$0.72\\pm 0.06$\\\\ \\sss Pierpaoli {\\etal} (2002)        &\\cite{Pierpaoli02}     &$0.77^{+0.05}_{-0.04}$\\\\ \\hline {\\bf Weak lensing:}\\\\ %95% limit, LCDM model, assuming Om=0.3 gamma=0.21,divide by two if using 1sigma \\sss Jarvis {\\etal} (2002)           &\\cite{Jarvis02}        &$0.71^{+0.06}_{-0.08}$\\\\ \\sss Brown {\\etal} (2002)            &\\cite{Brown02}         &$0.74\\pm 0.09$\\\\ \\sss Hoekstra {\\etal} (2002)         &\\cite{Hoekstra02a}     &$0.86^{+0.04}_{-0.05}$\\\\ \\sss Van Waerbeke{\\etal} (2002)      &\\cite{Waerbeke02}      &$0.97\\pm 0.06$\\\\ \\sss Bacon{\\etal} (2002)             &\\cite{Bacon02}         &$0.97\\pm 0.13$\\\\ \\sss Refregier{\\etal} (2002)         &\\cite{Refregier02}     &$0.94\\pm 0.14$\\\\ \\hline {\\bf CMB:}\\\\ \\sss Lahav {\\etal} 2001            &\\cite{Lahav01}         &$0.73\\pm 0.05$\\\\ \\sss Melchiorri \\& Silk 2002       &\\cite{Melch02a}        &$0.70\\pm 0.05$\\\\ \\sss Lewis {\\etal} 2002            &\\cite{Lewis02}         &$0.79\\pm 0.06$\\\\ \\sss This work           \t     &                       &$0.74\\pm 0.06$\\\\ % Bahcall et al 2002:  sigma_8*omega_m^0.6 = 0.35 \\hline \\end{tabular} } \\end{center} }   So is the CMB normalization too low or is the lensing normalization too high? As shown in Table 3, $\\sigma_8$ is emerging as the currently most controversial cosmological parameter, and it will be crucial to use further cosmological data to get to the bottom of this issue. % fly in the ointment. The table illustrates what has been emphasized by many authors, namely that the scatter between $\\sigma_8$ measurements using different techniques is substantially larger than the formal error bars, triggering unpleasant flashbacks to the bimodal and bitter debate about the Hubble parameter $h$. Although much work clearly needs to be done to resolve this issue, there are some encouraging indications that the $\\sigma_8$ gap may be closing. The CMB normalization has increased slightly in the last year, with the first acoustic peak growing by $10\\%$ between \\cite{xiaomin01} and our present compilation. CMB experiments can now be more effectively calibrated off of each other, and Archeops connects the calibration of DMR to that of the many experiments sensitive to the 1st acoustic peak and beyond. The cluster normalization has dropped somewhat because of recalibrations of the mass-temperature relation (see Table 3).  Although the lensing normalization has been uniformly high across groups \\cite{Bacon02,Hoekstra02a,Refregier02,Waerbeke02}, the largest and most recent data set released has now produced the much lower normalization $\\sigma_8\\approx 0.71$ \\cite{Jarvis02} (for $\\Om=0.3$). Indeed, the Hoekstra {\\etal} normalization may also be coming down by 5-10\\% (Hoekstra, private communication), partly because of a new and more accurate version of the Peacock \\& Dodds fitting formula \\cite{PeacockDodds96}.  Moreover, our analysis did not marginalize over the source redshift parameters $\\alpha=4.7$, $\\beta=1.7$ and $z_0=0.302$ from \\eq{SourceDistEq}, but merely used the best fit values from \\cite{Hoekstra02a}. Such marginalization would be expected to somewhat weaken the $\\sigma_8$ constraints that we obtained from the RCS lensing data. As an extreme example, we recomputed the shear power spectrum $P_\\kappa(\\l)$ replacing the characteristic source redshift parameter $z_s=0.302$ with the $3-\\sigma$ lower and upped limits on its value from \\cite{Hoekstra02a}. These extreme values $z_s=z=0.274$ and $z_s=0.337$ made the shear fluctuation amplitude $10-15\\%$ lower and higher, respectively, so lowering the source redshifts at the $2\\sigma$ level would close approximately half of the $\\sigma_8$ gap between the RCS and the CMB. Despite the current discord, there is thus real hope that a beautiful consistent picture of cosmology will eventually emerge.      \\subsection{Are the constraints too good to be true?}  In addition to the above-mentioned tension regarding the best fit values of some parameters, there is a slight puzzle regarding their error bars: some constraints seem a bit too good to be true. Specifically, the CMB error bars in Table 2 are in some cases comparable to those forecast for MAP \\cite{parameters2} using the Fisher matrix technique \\cite{karhunen}. This puzzling fact is not limited to the present paper, but appears rather generic. For instance, consider the scalar spectral index $\\ns$. The forecast for MAP (2 years of data including polarization) is accuracy $\\Delta\\ns=0.11$ with no prior assumptions, $\\Delta\\ns=0.045$ assuming $\\Ok=0$ and $\\Delta\\ns=0.021$ when adding SDSS galaxy clusterinng data (the final luminous red galaxy sample) \\cite{parameters2}. We found $\\Delta\\ns=0.06$ from current CMB alone (Table 2, assuming $\\Ok=0$), comparable to other studies in the recent literature. For instance, Lewis \\& Bridle \\cite{Lewis02} obtain $\\Delta\\ns=0.04$ including 2dF data, and the Boomerang team reports $\\Delta\\ns=0.06-0.08$ for various priors \\cite{Ruhl02}.  \\def\\D{{\\bf D}} \\def\\F{{\\bf F}} \\def\\Fcmb{\\F_{\\rm cmb}} \\def\\Flss{\\F_{\\rm 2df}} \\def\\SS{{\\bf\\Sigma}} \\def\\W{{\\bf W}}  To gain insight into this puzzle, we compute the Fisher matrix \\cite{karhunen} corresponding to the {\\it current} data. In the approximation that all the CMB information is coming from the mean rather than the covariance of the power, the equations in \\cite{karhunen} give \\beq{FisherEq} \\Fcmb = \\D^t\\W\\SS^{-1}\\W^t\\D, \\eeq where $\\SS$ is the $28\\times 28$ covariance matrix of our combined CMB power spectrum measurements from \\sec{CMBdata}, $\\W$ is the corresponding $\\lmax\\times 28$ window matrix and the $\\lmax\\times 28$ matrix \\beq{DdefEq} \\D_{\\l i}\\equiv {\\partial\\over\\partial p_i} \\delta T_\\l^2 \\eeq gives the derivatives of the power spectrum with respect to the cosmological parameters $p_i$. We compute an analogous Fisher matrix $\\Flss$ using the covariance matrix and window functions for the 2dF power spectrum from \\cite{Tegmark2df01}. We then evaluate the attainable error bars from CMB alone as $\\Delta p_i=(\\Fcmb^{-1})_{ii}^{1/2}$ and from CMB+2dF as $\\Delta p_i=([\\Fcmb+\\Flss]^{-1})_{ii}^{1/2}$. For the scalar spectral index with an $\\Ok=0$ prior, this gives $\\Delta\\ns=0.16$ from CMB alone and $\\Delta\\ns=0.13$ for from CMB+2dF, a factor 2-3 larger than the above-mentioned likelihood analysis results from us and others.  Comparing with our Fisher matrix results, the most suspicious error bar in Table 2 is that on the optical depth $\\tau$, where we quoted $\\Delta\\tau=0.06$ from CMB alone and the Fisher matrix gives $\\Delta\\tau=0.18$. For the $\\tau$ case, there is wide scatter in the literature. On the optimistic side, the CBI team reported $\\tau=0.09^{+.14}_{-0.07} $\\cite{CBI02a}, and a strong detection of reionization was claimed from earlier data \\cite{Schmalzing00}. % Schmalzing et al (2000) gave tau=?_/-? from Maxima + lots of priors) (zion>15), astro-ph/0010063 On the pessimistic side, the 2dF and Boomerang teams report a $2\\sigma$ limit $\\tau<0.5$ \\cite{Efstathiou01,Ruhl02}, % with most recent studies intermediate with results like $\\tau\\simlt 0.2$. % New Ruhl et al give tau < 0.51 (95%), <0.2 (1 sigma), <0.3 (95%, flat prior) % Melchiorri \\& Oedman get tau<.25 (68\\%) from CMB alone, incl Archeops (Ok=0). in line with the Fisher forecast.  If our $\\tau$ errors are indeed too small due to some artifact of our method, this would weaken the constraints on the normalization $\\sigma_8$, since CMB constrains mainly the combination $\\sigma_8 e^{-\\tau}$. However, the conclusion that $\\sigma_8=0.9$ forces the uncomfortably high value $\\tau\\approx 0.2$ would still stand, since it follows directly from this scaling relation.  % Boomerang\\cite{0104460} gives tau<0.2 with Ok=0+weak priors % (where red curve drops below 0.6 in their Fig 6) %CBI team 0205387: %All data  0.16+.18/-.13 %All data  0.09+.14/-.07 %Wang {\\etal} \\cite{Xiaomin01} obtain $\\tau<0.18$ ($2\\sigma$). % Lewis et al claim neutrino mass sum m_nu<1 eV  Efstathiou \\& Bond \\cite{EfstathiouBond98} have shown that there are physically understandable reasons why the Fisher matrix approach can significantly overestimate the errors on cosmological parameters, especially when strong degeneracies are present and that the curvature of a banana-shaped allowed region becomes important. This is indeed the case with $\\tau$, with its strong degeneracy with the normalization parameter $\\As$, and a rigorous Bayesean confidence interval taking into account the asymmetry $\\tau\\ge 0$ (neglected in the Fisher treatment) also matters. ``Weak'' priors can also have a strong effect, notably on $h$ and gravity waves. Hopefully, our concerns will therefore go away when the MAP data arrives, in particular when all major degeneracies are broken with other datasets and priors. Since \\eq{FisherEq} is trivial to apply, it would be prudent to include a Fisher analysis of the data used in all future parameter constraint papers, as a reality check.    \\subsection{Outlook}  In conclusion, the pace at which new measurements have tightened up constraints on cosmological parameters is breathtaking. The constraints become particularly tight if one uses theoretical prejudice to impose spatial flatless $\\Ok=0$ like we have done, interpreting the CMB indication $\\Ok\\approx 0$ as support for inflation and perfect flatness and thereby breaking the worst parameter degeneracy of all by fiat. It can be argued that precision cosmology is already here, before MAP, and now faces its first baptism of fire: will the current precision measurements of cosmological parameters agree with MAP? Next month we will know whether the whole edifice goes down in flames or stands stronger than ever.     The authors wish to thank Gary Bernstein, Karim Benabed and Roman Scoccimarro for helpful comments. This work was supported by NSF grants AST-0071213, AST-0134999, AST-0098606 and PHY-0116590, NASA grants NAG5-9194, NAG5-11099, NAG5-10923 and NAG5-10924 and a Keck foundation grant. MT and MZ are David and Lucile Packard Foundation Fellows MT is a Cottrell Scholar of Research Corporation.   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%% REFERENCES: %%%%%%%%%%%%%%%%%%%%%%%%%  %\\clearpage %\\end{multicols}  %\\vskip-1.0cm"
    },
    "0212/astro-ph0212517_arXiv.txt": {
        "abstract": "We report an investigation of cosmological parameters based on the measurements of anisotropy in the cosmic microwave background radiation (CMB) made by \\acbar.  We use the \\acbarsp data in concert with other recent CMB measurements to derive Bayesian estimates of parameters in inflation-motivated adiabatic cold dark matter models. We apply a series of additional cosmological constraints on the shape and amplitude of the density power spectrum, the Hubble parameter and from supernovae to further refine our parameter estimates. Previous estimates of parameters are confirmed, with sensitive measurements of the power spectrum now ranging from $\\ell \\sim 3$ to $2800$. Comparing individual best model fits, we find that the addition of $\\Omega_{\\Lambda}$ as a parameter dramatically improves the fits. We also use the high-$\\ell$ data of \\acbar, along with similar data from CBI and BIMA, to investigate potential secondary anisotropies from the Sunyaev-Zeldovich effect.  We show that the results from the three experiments are consistent under this interpretation, and use the data, combined and individually, to estimate $\\sigma_8$ from the Sunyaev-Zeldovich component. ",
        "introduction": "\\begin{split} \\ln{\\mathcal L}({\\bf{\\mathcal C}}) = &\\ln{\\mathcal L}(\\overline{{\\bf\\mathcal C}}) -\\\\ &\\frac{1}{2}\\sum\\limits_{BB^{\\prime}}(Z_B - \\overline{Z}_B) {\\mathcal F}_{BB^{\\prime}}^{(z)}(Z_{B^{\\prime}} - \\overline{Z}_{B^{\\prime}}) \\ , \\end{split}\\end{equation} where $\\overline{Z}_B =\\ln{(\\overline{{\\mathcal C}}_B + x_B)}$ is the value of the lognormal parameter at the position of maximum likelihood $\\overline{{\\mathcal C}}_B$ and ${\\mathcal F}_{BB^{\\prime}}^{(z)}$ is the curvature matrix transformed into the lognormal variables, \\begin{equation} {\\mathcal F}_{BB^{\\prime}}^{(z)} = (\\overline{{\\mathcal C}}_B + x_B) {\\mathcal F}_{BB^{\\prime}}^{({\\mathcal C})} (\\overline{{\\mathcal C}}_B^{\\prime} + x_B^{\\prime}) \\ . \\end{equation} Provided with the maximum likelihood band powers $\\overline{{\\mathcal C}}_B$, the lognormal offsets $x_B$, the curvature matrix ${\\mathcal F}_{BB^{\\prime}}^{({\\mathcal C})}$, and the window functions $\\varphi_{B\\ell}$, the likelihood of the parameter set $\\vec{y}$, given the data, can be computed.  Our set $\\vec{y}$ consists of seven cosmological parameters: {$\\Omega_k=(1-\\Omega_{tot}), \\Omega_{\\Lambda}, \\omega_{cdm}, \\omega_b, n_s, \\tau_C, \\ln{{\\mathcal C}_{10}}$}.  The total energy density of the universe in units of critical density $\\Omega_{tot}$ is linked to the global curvature of space:  negatively curved for $\\Omega_{tot} < 1$, positively curved for $\\Omega_{tot} > 1$ and flat for $\\Omega_{tot} = 1$.  The total energy density has 3 constituents: vacuum ($\\Omega_\\Lambda$), matter and relativistic particles.  The relativistic energy density is currently negligible.  The matter density is split into two types, baryonic matter ($\\Omega_b \\equiv \\omega_b/h^2$), which interacts with electromagnetic radiation, and cold dark matter ($\\Omega_{cdm} \\equiv \\omega_{cdm}/h^2$),  which does not. The total matter density is denoted $\\Omega_m = \\Omega_b + \\Omega_{cdm}$. The amplitude of the CMB power spectrum at $\\ell = 10$, $\\ln{{\\mathcal C}_{10}}$ gives the overall amplitude of the primordial fluctuations. This quantity is well-constrained by the COBE-DMR observations~\\citep{bennett96}. The full COBE-DMR power spectrum as described in~\\citet{bond00} is included in all of our parameter analyses. The spectral index of primordial density perturbations, $n_s$, parameterizes the variation in the fluctuation power as a function of length scale; $n_s = 1$ corresponds to scale invariance.  The universe reionized at some point between decoupling and the present.  After reionization, CMB photons scatter further. $\\tau_C$ is the Compton optical depth (from decoupling to present) due to such scattering.  High $\\tau_C$ diminishes CMB power by a factor of $\\exp[-2\\tau_C]$ over most of the $\\ell$ range, though not in the DMR range.  Many more parameters than our basic 7 may be needed to completely describe inflationary models. These include the gravity-wave induced tensor amplitude and tilt, variations of tilt with wavenumber, relativistic particle densities, more complex dynamics associated with the dark energy $\\Omega_\\Lambda$, etc.  For example, a gravity-wave induced component is expected in the largest class of inflation models, and has impact on the spectrum in the $\\ell < 100$ region. Although far from the region where \\acbarsp has its impact, it can affect the amplitude, tilt, and Compton depth. If one constrains the tilts and amplitudes of the tensor component to those motivated by simple inflation models, there is little change in the other parameters.  It is possible that secondary anisotropies, such as the Sunyaev-Zeldovich effect investigated later in this paper, could contribute significant power to the highest $\\ell$ band of our measurement.   However, the detection of anisotropy in that band is only $1.1\\sigma$, or $0.9\\sigma$ above the best-fit model primary CMB angular power spectrum.  Thus, for the purpose of cosmological parameter estimation from the primary CMB signal, we can safely ignore the effects of potential SZE contamination.  To derive estimates of cosmological parameters, we compare our data with the primary CMB power spectra predicted by combinations of cosmological parameters $\\vec{y}$. In this comparison, we vary $\\ln{{\\mathcal C}_{10}}$ continuously, while the other parameters take the discrete values listed in Table~\\ref{paramlist}. An angular power spectrum is generated for each set of discrete parameter values, forming a grid. Given a CMB dataset (\\eg \\acbarsp or a combination of various measured power spectra), the likelihood is then calculated for each point on the seven dimensional grid.  To compute the likelihood that a particular parameter $X$ has a value $x_0$, the seven dimensional grid of likelihoods is integrated over the other six parameters, holding the parameter of interest fixed at $x_0$. This method, known as marginalization, involves calculating \\begin{equation}\\label{prior} L(X=x_0) = \\int \\delta(X-x_0)P_{prior}(\\vec{y}){\\mathcal L}(\\vec{y})d\\vec{y} \\ , \\end{equation} where $\\delta(x)$ is the usual delta function and the prior $P_{prior}(\\vec{y})$ is discussed below.  For each model on the parameter grid, along with the overall amplitude parameter ${\\cal C}_{10}$, we continuously vary the beam widths $\\sigma_{bi}$ and calibrations $A_i$ for each experiment $i$ about their estimated values $\\overline{\\sigma_{bi}}$ and unity, to take into account the uncertainties in the respective measurements.  We approximate the beam-uncertainty and calibration-uncertainty ``prior'' probabilities by Gaussians in $\\Delta (\\sigma_{bi})^2$ and $\\ln A_i^2$, respectively. The modification to the bandpower as a result of the uncertainty $\\Delta (\\sigma_{bi})^2$ is modeled by $\\exp[-\\langle(\\ell+\\frac{1}{2})^2\\rangle_B \\Delta (\\sigma_{bi})^2]$, with $\\langle(\\ell+\\frac{1}{2})^2\\rangle_B = {\\mathcal I}(\\varphi_{B\\ell}(\\ell+\\frac{1}{2})^2)/ {\\mathcal I}(\\varphi_{B\\ell}) $. The overall impact of the calibration and beam uncertainties is that the combination $\\ln {\\cal C}_{10} + \\sum_i \\ln A_i^2 - \\langle(\\ell+\\frac{1}{2})^2\\rangle_B \\Delta (\\sigma_{bi})^2$ is adjusted for each grid parameter combination to give the best fit with errors. Marginalization over the continuous parameters is done by calculating a Fisher matrix, and assuming a Gaussian distribution in the posterior distributions in the $\\{ \\ln {\\cal C}_{10}, \\ln A_i^2, \\Delta (\\sigma_{bi})^2\\} $ variables. Marginalization over the grid parameters is done by discrete integration.  In addition to the parameters given in Table~\\ref{paramlist}, we examine four ``derived'' parameters, which are functions of our basic 7 variables, $\\{t_0$, $h$, $\\sigma_8^2$, and $\\Gamma_{\\rm eff} \\}$. Here $h$ is the value of the Hubble expansion parameter, $H_0 = 100 ~h$ km/s/Mpc and is given by \\begin{equation} h = {\\sqrt \\frac{\\omega_b + \\omega_{cdm}}{1 - \\Omega_k - \\Omega_\\Lambda}} \\ . \\end{equation} The age of the universe is \\begin{equation} t_0({\\rm Gyr}) \\approx \\frac{9.778}{h} \\int_0^1 \\frac{2x^2}{\\sqrt{\\Omega_m+\\Omega_kx^2+\\Omega_{\\Lambda}x^3}}dx \\ , \\end{equation} where we have dropped the minor effects of relativistic particles here, but not when we do our actual comparisons. The variance in the (linear) density fluctuation spectrum on the scale of clusters of galaxies ($8 h^{-1}$ Mpc) is $\\sigma_8^2$. The shape of the linear density power spectrum is described by the parameter \\begin{equation}\\begin{split} &\\Gamma_{\\rm eff}=\\Gamma + (n_s- 1)/2 ,\\quad {\\rm where}\\\\ &\\Gamma = \\Omega_{m} \\, h \\, \\, e^{-(\\Omega_b(1+\\Omega_{m}^{-1}\\sqrt{2h}) -0.06)} \\ . \\end{split}\\end{equation} Note that $\\Gamma_{\\rm eff} \\approx \\Omega_m h$ with corrections due to baryon density and spectral tilt $n_s$ over the region probed by large scale structure observations that approximate the dominant dependences. The derived parameters are not marginalized over, and do not define dimensions in the parameter grid because they are determined given a set of parameters $\\vec{y}$. The derived parameters are calculated at each point on the model grid.  We can make Bayesian estimates of those parameters using equation~\\eqref{prior} with $X$ being the derived parameter of interest.  We can also use the derived parameters to account for other cosmological information.  Estimates of these parameters from other, non-CMB, observations can be included as ``prior constraints\". Equation~\\eqref{prior} is cast in a form that gives a likelihood as a function of our parameters, and priors $P_{prior}(\\vec{y})$.  All our analyses include the loose prior implicitly imposed by the edges of the database given in Table~\\ref{paramlist}, and two other very weak priors that are generally accepted by most cosmologists:  models are restricted to those where the current age of the universe is $t_0 > 10$Gyr, and where the matter density $\\Omega_m$ is greater than 0.1.  We add to these constraints a series of additional priors: \\begin{itemize}  \\item {\\bf weak-h :} $0.45 < h < 0.90$. This is a tophat restriction, designed to allow the CMB data, rather than arguable priors, to drive the results.  \\item {\\bf LSS :}  We employ two constraints based on large scale structure (LSS) observations, on the combinations $\\sigma_8 \\Omega_m^{0.56}$ and $\\Gamma_{\\rm eff}$. For both priors, we adopt a Gaussian distribution convolved with a top-hat distribution, characterized by the parameters $\\sigma_8 \\Omega_m^{0.56} = 0.47_{-.02, -.08}^{+.02, +.11}$ and $\\Gamma_{\\rm eff} = 0.21_{-.03, -.08}^{+.03, +.08}$, where the first error gives the 1-$\\sigma$ point of the Gaussian distribution, and the second error gives the extent of the tophat distribution. Our basic philosophy is to adopt priors that are not overly restrictive since the LSS data is still improving. The motivation for the choice and the discussion of the LSS data is given in \\citet{bond02b}. The $\\Gamma_{\\rm eff}$ distribution encompasses recent results from the 2dF and SDSS surveys.  The $\\sigma_8$ distribution encompasses results from recent weak lensing surveys. It also covers many of the cluster abundance determinations using X-ray temperature and other cluster data.  \\item {\\bf LSS($low$--$\\sigma_8$):} There are currently a few cluster abundance estimations that point to values of $\\sigma_8$ that are lower than the weak lensing estimates. Although our standard LSS prior takes most of these variations into account by its spread, we have also tested the effect of shifting the entire $\\sigma_8$-distribution downward by 15\\%, to $\\sigma_8 \\Omega_m^{0.56} = 0.40_{-.02, -.08}^{+.02, +.11}$. We keep the $\\Gamma$ prior the same.  \\item {\\bf HST-h :}  We strengthen our $h$ prior, based on the Hubble Space Telescope Key Project measurement~\\citep{freedman00} of the Hubble constant, $h = 0.72 \\pm 0.08$. This is a Gaussian prior with the stated error as the 1--$\\sigma$ points.  \\item {\\bf Strong data :}  We combine the LSS prior given above with the HST-$h$ prior, and add a constraint (in the $\\Omega_{\\Lambda}$ vs. $\\Omega_{tot}$ plane) based on surveys of the brightness vs. redshift relation of Type 1a Supernovae~\\citep{perlmutter99,riess98}.  \\item {\\bf Flat :}  Inflation models generally predict a flat geometry, and recent evidence supports this~\\cite[eg][]{debernardis00,halverson01}.  For the converted, we investigate the effects of holding $\\Omega_{tot}$ equal to one. \\end{itemize} ",
        "conclusions": "The \\acbarsp data, the most sensitive to date in the damping-tail region, are in good agreement with predictions of flat $\\Lambda$-CDM models with adiabatic initial perturbations. Considering the \\acbarsp data together with other recent CMB results, we find that the addition of a single parameter to the model, $\\Omega_\\Lambda$, dramatically improves the best fit, with an improvement of $\\Delta \\chi^2 = 20$ upon adding that one parameter. Using very weak cosmological priors ($\\Omega_m > 0.1$, age $> 10$Gyr, $0.45 < h < 0.90$), the 3-$\\sigma$ lower limit on the cosmological constant rises to $\\Omega_\\Lambda > 0.136$ upon including the \\acbarsp data.   We find that the addition of \\acbarsp data to the current CMB set does not lead to substantial improvements in the 1-$\\sigma$ estimates of the canonical cosmological parameters.  However, in an eigenmode analysis, the addition of \\acbarsp data does improve the rotated parameter uncertainties, indicating that in this case the lack of improved errors on the pure cosmological parameters is probably dominated by degeneracies between those parameters.  We fit a SZE component to the data using \\acbarsp and other measurements at high-$\\ell$.  Our estimates for the value of the effective quantity $\\sigma_8^{\\rm SZ}=(\\Omega_b h/0.035)^{0.29}\\sigma_8$ show an improvement over previous estimates using only CBI and BIMA observations. Although CBI and \\acbarsp alone do not provide lower bounds on $\\sigma_8^{\\rm SZ}$, the combination of the two observations results in a detection greater than 3-$\\sigma$ which is independent of the BIMA lower bound.  Our fits show that at a fiducial value of $\\Omega_b h=0.035$ the central values for $\\sigma_8$ are consistently higher than other estimates obtained using cluster data, weak lensing surveys and primary CMB observations (see \\citet{bond02b} for a recent survey of $\\sigma_8$ estimates). However, the results overlap with most other estimates at the 2-$\\sigma$ level. As an indication, our LSS prior, a smoothed tophat on the combination $\\sigma_8\\Omega_m^{0.56}$, translates roughly into a smoothed tophat of $\\sigma_8=0.92^{+0.21}_{-0.15}$ at $\\Omega_m=0.3$. Any {\\it statistical} inconsistency therefore, appears to be mild.  Furthermore, {\\it systematic} uncertainties in the estimates have not yet been taken into account.  The difference displayed by the numerical and analytical based results of about 10\\% is indicative of the agreement between the two methods for predicting the SZE power spectrum \\citep{komatsu02wc}. In addition, entropy injection may have a significant effect on the SZE power spectrum. These effects would change the shape of the SZE template and impact directly on our determination of $\\sigma_8$. Nevertheless, we conclude that the in the context of the phenomenological models adopted in this work, the data is consistent with a SZE component at $\\sigma_8$ values near the high end of independent estimates."
    },
    "0212/astro-ph0212503_arXiv.txt": {
        "abstract": "We study the dynamical evolution of a stationary, axisymmetric, and perfectly conducting cold accretion disk containing a large-scale magnetic field around a Kerr black hole, trying to understand the relation between accretion and the transportation of angular momentum and energy. A one-dimensional radial momentum equation is derived near the equatorial plane, which has one intrinsic singularity at the fast critical point. We solve the radial momentum equation for solutions corresponding to an accretion flow that starts from a subsonic state at infinity, smoothly passes the fast critical point, then supersonically falls into the horizon of the black hole. The solutions always have the following features: 1)~The specific energy of fluid particles remains constant but the specific angular momentum is effectively removed by the magnetic field. 2)~At large radii, where the disk motion is dominantly rotational, the energy density of the magnetic field is equipartitioned with the rotational energy density of the disk. 3)~Inside the fast critical point, where radial motion becomes important, the ratio of the electromagnetic energy density to the kinetic energy density drops quickly. The results indicate that: 1)~Disk accretion does not necessarily imply energy dissipation since magnetic fields do not have to transport or dissipate a lot of energy as they effectively transport angular momentum. 2)~When resistivity is small, the large-scale magnetic field is amplified by the shearing rotation of the disk until the magnetic energy density is equipartitioned with the rotational energy density, ending up with a geometrically thick disk. This is in contrast with the evolution of small-scale magnetic fields where if the resistivity is nonzero the magnetic energy density is likely to be equipartitioned with the kinetic energy density associated with local random motions (e.g., turbulence), making a thin Keplerian disk possible. ",
        "introduction": "\\label{sec1}  Accretion onto a gravitating object is believed to be a principal energy source for powering many astrophysical systems, including active galactic nuclei (AGN), quasars, and some types of stellar binaries \\cite{fra02}. Since accreting gases far from the central object typically have sufficiently large specific angular momentum compared to a particle moving on a circular orbit around and near the central object, the formation of an accretion disk is unavoidable. To accrete onto the central object and release gravitational energy, gases at large distance must lose a lot of angular momentum. It is usually assumed that some kind of viscosity slowly operates in a disk to transport angular momentum outward and dissipate energy, so that to a very good approximation in the disk region particles move on circular orbits with a small radial velocity superposed.  The simplest stationary and axisymmetric model of accretion disk is a geometrically thin Keplerian disk, where it is assumed that in the disk region the gravity of the central object is approximately balanced by the centrifugal force so that disk particles move approximately on Keplerian circular orbits \\cite{sha73,nov73,pag74}. Then, at each radius, the specific angular momentum and the specific energy of disk particles are given by that of a test particle moving on the Keplerian circular orbit at that radius. When the central object is a black hole, the Keplerian disk has an inner boundary at the marginally stable circular orbit, inside which the gravity is too strong to be balanced by the centrifugal force \\cite{bar72,nov73,pag74}. The total efficiency of converting mass into energy is defined by \\begin{eqnarray} \\varepsilon \\equiv \\frac{\\delta {\\cal E}}{\\delta M} \\;, \\end{eqnarray} where $\\delta {\\cal E}$ is the total released energy during the process of accretion, $\\delta M$ is the corresponding accreted mass \\footnote{Throughout the paper we use geometrized units $G = c = 1$, where $G$ is the Newtonian gravitational constant, $c$ is the speed of light.}. Then, the total efficiency of a Keplerian disk around a Kerr black hole is given by the difference between the specific energy at infinity ($E_\\infty = 1$) and the specific energy on the marginally stable circular orbit ($E_{\\rm ms}$) \\begin{eqnarray} \\varepsilon_{\\rm Keplerian} = 1 - E_{\\rm ms} \\;. \\label{ek} \\end{eqnarray}  Thin Keplerian disks correspond to the most efficient case of releasing the gravitational energy of an accretion flow. For a thin Keplerian disk around a Schwarzschild black hole, the total efficiency is $\\varepsilon_{\\rm Keplerian} \\approx 0.06$. For a thin Keplerian disk around a ``canonical'' Kerr black hole of specific spinning angular momentum $a = 0.998 M$, where $M$ is the mass of the black hole, the total efficiency is $\\varepsilon_{\\rm Keplerian} \\approx 0.32$ when the effect of radiation capture is considered \\cite{tho74}. These efficiencies of converting mass into energy are indeed much higher than nuclear fusion, the latter is $\\approx 0.007$ for hydrogen. The released gravitational energy in a thin Keplerian disk is assumed to be efficiently dissipated by the internal viscosity and instantly radiated away from the surfaces of the disk.  Though the model of geometrically thin and optically thick Keplerian accretion disks is supported by the appearance of Big Blue Bump in the spectra of all bright AGN (\\cite{shi78,kor99,kis03} and references therein), it is challenged by the observational fact that most nearby galactic nuclei are much less active---sometimes not active at all (\\cite{nar02} and references therein). These objects have luminosities that are lower than what a thin Keplerian disk model usually predicts by many orders of magnitudes. In order to interpret this phenomenon, different disk models are proposed where either the radiative efficiency is assumed to be low so that the released gravitational energy is converted into entropy and internal energy trapped in the flow and carried into the central black hole by the flow, or the mass accretion rate is assumed to be low and most of the released gravitational energy is carried away by a mass outflow \\cite{nar94,bla99b,sto99,nar00,qua00,igu03}. These alternative models are summarized by Blandford \\cite{bla99a} and Narayan \\cite{nar02}. Though they predict a different disk luminosity, these models are in common with the thin Keplerian disk in the following aspect: the mechanical energy (= gravitational potential energy + kinetic energy) of the disk flow is efficiently converted into other forms of energy.  In the theory aspect, the proposed disk viscosity is poorly understood. Investigations on the magnetohydrodynamics (MHD) in a differentially rotating disk suggest that magnetic fields may play the role of disk viscosity (\\cite{bal98,bla02} and references therein). Due to the existence of magnetorotational instability, MHD turbulence is easily generated in a differentially rotating disk, which effectively transports angular momentum outward and drives accretion (\\cite{bal91,sto96,haw96,haw00,haw02} and references therein). However, very few models of disk radiation are based on these ideas, and generally angular momentum transport is treated through a parameterized viscous stress---the so-called $\\alpha$-viscosity following Shakura \\& Sunyaev \\cite{sha73,kor99}. This is because of the fact that though it is well established that magnetic fields effectively transport angular momentum and drive accretion, it is far from clear how magnetic fields transport and dissipate energy in the meantime. Numerical simulations are a powerful tool for studying the nonlinear evolution of magnetic fields. However, numerical simulations usually have large numerical dissipation due to finite spatial resolutions, which makes energy transportation and dissipation cannot be accurately studied with current numerical simulations \\cite{haw95,ryu95,bal98,bla02}. Therefore, our current understanding on how magnetic fields transport and dissipate energy is much worse than our understanding on how magnetic fields transport angular momentum.  In this paper we ask and try to answer the following question: {\\it Do magnetic fields transport or dissipate energy as efficiently as they transport angular momentum?} Or, more precisely, {\\it Do magnetic fields have to transport or dissipate a lot of energy as they efficiently transport angular momentum in a disk accretion flow?}  To answer this question (or, to get some insight into the answer), we use an analytical model to study the dynamical effects of magnetic fields in an accretion disk around a Kerr black hole. The disk is assumed to be stationary, axisymmetric, perfectly conducting, and contain a large-scale magnetic field~\\footnote{A large-scale magnetic field in an accretion disk can arise in at least two ways: 1)~Inherited from the progenitor of the disk. For example, merger of a magnetized neutron star with a black hole will naturally lead to the formation of a disk with a large-scale magnetic field. 2)~Generated from small-scale magnetic fields through reconnection. See, e.g., C. A. Tout \\& J. E. Pringle, Mon. Not. R. Astron. Soc. {\\bf 281}, 219 (1996).}. Both the magnetic field and the velocity field of the disk fluid have only radial and azimuthal components in a small neighborhood of the disk central plane $\\theta = \\pi/2$, where $\\theta$ is the polar angle from the symmetry axis of the black hole. In other words, near the central plane of the disk, each magnetic field line lies on a surface of constant $\\theta$, and the same is true for each fluid stream line. (Similar approach is usually used to study the inflow/outflow problem associated with a low density magnetosphere around a black hole or neutron star, see \\cite{cam86,cam86a,tak90,hir92,dai02,tak02}, and \\cite{pun01} for a review.) For simplicity, we neglect the thermal pressure in the disk by assuming that the accretion flow is cold. Then, making use of the conservation of rest mass, energy, and angular momentum, Maxwell's equations and dynamical equations are reduced to a set of one-dimensional equations with radius $r$ as the only variable. We will self-consistently solve these equations, to look for the relation between accretion and the transportation of angular momentum and energy.  Gammie \\cite{gam99} has used the model to study the flow on the equatorial plane in the plunging region between the inner boundary of the disk and the horizon of the black hole. Gammie's model is further explored by Li \\cite{li03} and tested with numerical simulations by De Villiers \\& Hawley \\cite{vil02}. In this paper we extend the model into the disk region, and into a small neighborhood of the equatorial plane. This extension allows us to investigate a disk accretion flow driven by magnetic fields from a global view: from the subsonic far outer disk region down to the supersonic deep plunging region~\\footnote{In this paper we use the word ``sonic'' to refer to ``magnetosonic'' since we have assumed that the gas pressure is zero.}. We emphasize that such a model can only be applied to a large-scale and ordered magnetic field. For small-scale and chaotic magnetic fields, the gradient of the magnetic field is important, making the solutions invalid when even a small nonzero resistivity is introduced (see Sec.~\\ref{sec6} for detail).  The paper is organized as follows: In Sec.~\\ref{sec2} we exactly solve Maxwell's equations for the model outlined above. In Sec.~\\ref{sec3} we write down the equations for the conservation of rest mass, the conservation of angular momentum, and the conservation of energy. We show that the solutions always correspond to an accretion flow with {\\it constant specific energy} when reasonable boundary conditions at infinity are satisfied. Then we derive a one-dimensional radial momentum equation corresponding to the magnetized disk accretion flow.  In Sec.~\\ref{sec4} we discuss the analytical properties of the radial momentum equation, the solutions on the horizon of the black hole, and the asymptotic solutions at infinity. We show that the fast critical point is the only intrinsic singularity in the radial momentum equation. In Sec.~\\ref{sec5} we present some explicit solutions representing a magnetized disk flow with constant specific energy accreting from infinity onto a central black hole. In Sec.~\\ref{sec6} we discuss the effect of finite resistivity and compare the evolution of large-scale magnetic fields to that of small-scale magnetic fields. In Sec.~\\ref{sec7} we summarize the results and draw our conclusions.  In Appendix~\\ref{appa} we make some effort to generalize our results to more general and more complicated models. In Appendix~\\ref{appb} we present the Newtonian limit of our solutions. ",
        "conclusions": "\\label{sec7}  With an analytical model we have investigated the dynamics of a cold accretion disk containing a large-scale magnetic field around a Kerr black hole, trying to understand the process of angular momentum and energy transportation. A one-dimensional radial momentum equation is derived near the equatorial plane of the black hole, which has one intrinsic singularity at the fast critical point. Any physical solution corresponding to an accretion flow starting from infinity subsonically must smoothly pass the fast critical point in order to reach the horizon of the black hole. Because of this requirement, among the four integral constants ($c_0$, $\\Omega_\\Psi$, $f_{\\rm L}$, and $f_{\\rm E}$) only three are independent: one of them has to be determined as a function of the other three.  In order for $r^2 B^r$ and $B^2/4\\pi\\rho_{\\rm m}$ to be finite at infinity, the constant $\\Omega_\\Psi$ has to be zero (Secs.~\\ref{sec2} and \\ref{sec3.1}). This indicates that in the stationary state the magnetic field and the velocity field of the fluid must be parallel to each other. When this condition is satisfied ($\\Omega_\\Psi= 0$), the specific energy of the fluid particles remains constant during the motion, but the specific angular momentum changes.  The radial momentum equation has been solved with the assumptions that $\\Omega_\\Psi = 0$, $c_0^2/r_{\\rm g}^2 \\ll 1$, and $f_{\\rm E} = -1$ (corresponding to an accretion flow with the outer boundary at infinity). The results can be summarized as follows: \\begin{itemize} \\item[1.]{The accretion flow starts from a subsonic state at large radii, passes the fast critical point near the marginally bound circular orbit, then enters the black hole supersonically (see, e.g., Figs.~\\ref{fig4}).} The solutions are super-Keplerian (Figs.~\\ref{fig5}, \\ref{fig8}, \\ref{figk3} and \\ref{figk4} upper panels), whose asymptotic behaviors on the black hole horizon and at infinity are summarized in Table~\\ref{tab1}. \\item[2.]{The specific energy of the fluid particles remains constant, but the specific angular momentum is effectively removed by the magnetic field and transported outward (Figs.~\\ref{fig5}, \\ref{figk3} and \\ref{figk4} upper panels). The rate of transporting angular momentum is similar to that for a thin Keplerian disk, though the specific angular momentum that the flow finally carries into the black hole is somewhat larger.} \\item[3.]{At large radii the electromagnetic energy density is about equipartitioned with the kinetic energy density of the fluid particles, but near and inside the fast critical point the ratio of the electromagnetic energy density to the kinetic energy density is $\\ll 1$ (Figs.~\\ref{fig7}, \\ref{figk3} and \\ref{figk4} lower panels). This is due to the fact that at large radii the velocity of the fluid is predominantly toroidal so the magnetic field is sufficiently amplified by the shear rotation, while near and inside the fast critical point the radial velocity of the fluid becomes important then the magnetic field has no time to get sufficient amplification (see Fig.~\\ref{fig6}).} \\item[4.]{When the black hole is Schwarzschild, the electric field measured by a local static observer is zero and the ratio of the magnetic energy density to the kinetic energy density approaches zero on the horizon of the black hole (Fig.~\\ref{fig7}).} \\item[5.]{When the black hole is Kerr, due to the frame dragging effect the electric field measured by a LNRF observer is nonzero. The electric field decays rapidly with increasing radius, but near the horizon the electric field makes a significant contribution to the total electromagnetic energy density. On the horizon, the ratio of the total electromagnetic energy density to the kinetic energy density, which is contributed equally by electric field and magnetic field, is tiny but nonzero (see Figs.~\\ref{figk3} and \\ref{figk4}}, lower panels). \\item[6.]{The specific angular momentum of the fluid particles as they reach the horizon, $L_{\\rm H}$, is a strong function of $a/M$ and $c_0/r_{\\rm g}$, though in the limit $c_0^2/r_{\\rm g}^2 \\ll 1$ $L_{\\rm H}$ becomes insensitive to the change in $c_0/r_{\\rm g}$ (Fig.~\\ref{figk5}). For a positive $a$, $L_{\\rm H}$ is always positive. For a zero or negative $a$, $L_{\\rm H}$ becomes zero or negative for sufficiently large $c_0/r_{\\rm g}$.} \\end{itemize}  With the existence of solutions with constant specific energy for a magnetized accretion flow, the question posed in the Introduction ({\\it Do magnetic fields transport or dissipate energy as efficiently as they transport angular momentum?}) is at least partly answered: {\\it Magnetic fields do not have to transport or dissipate a lot of energy as they efficiently transport angular momentum.} This has important implications to the theory of accretion disk: an accretion disk driven by magnetic fields may have a very low radiation efficiency. People generally believe that an accretion disk usually has a radiation efficiency much higher than a spherical accretion flow. The belief comes from the following argument: a disk has a lot of angular momentum which must be removed in order for the disk material to be able to accrete onto the central compact object. In the meantime, a lot of energy is expected to be removed also. The results in this paper show that the specific energy of fluid particles does not have to change when its angular momentum is being removed. In fact the solutions with constant specific energy are the unique solutions for the model studied in this paper, all other solutions have bad asymptotic behaviors at infinity (namely, $r^2 B^r$ and $B^2/4\\pi\\rho_{\\rm m}$ diverge), if the specific angular momentum of the fluid is unbounded at infinity.  We have ignored energy dissipation by assuming that the disk plasma is perfectly conducting. Then, inevitably, at large radii the magnetic energy density is equipartitioned with the kinetic energy density, indicating that the accretion flow must be geometrically thick. The solutions presented in the text apply only to the region near the central plane of the disk. However, in App.~\\ref{appa} we show that the solutions can be easily generalized to more general and more complicated models, including the disk region well above or well below the equatorial plane. So, solutions with constant specific energy exist for more general and more complicated models.  The effects of finite resistivity are discussed in Sec.~\\ref{sec6}. When the resistivity is nonzero, an accretion disk driven by a large-scale and ordered magnetic field behaves very differently from that driven by small-scale and chaotic magnetic fields. For a large-scale and ordered magnetic field, in a stationary and axisymmetric state the magnetic energy density is approximately equipartitioned with the rotational energy density at large radii if the resistivity is small. For small-scale and chaotic magnetic fields that reverse directions frequently, even if the resistivity is small, in the equilibrium state the magnetic energy density is likely to be equipartitioned with the kinetic energy density associated with local random motions (e.g., turbulence).  A finite resistivity in the disk will give rise to a nonzero radiation efficiency. However, a nonzero radiation efficiency may also arise in the following way: a disk corresponding to our solutions has a finite outer radius $r_0$ and beyond which it is joined to a thin Keplerian disk. Then, the arguments in Sec.~\\ref{sec3.1} indicate that the region inside $r_0$ might have a radiation efficiency $\\varepsilon_{\\rm in} \\alt \\Omega_0 L_0 \\sim r_{\\rm g}/r_0\\,$, assuming that $r_0\\gg r_{\\rm g}\\,$. The thin Keplerian disk outside $r_0$ would give rise to a radiation efficiency $\\varepsilon_{\\rm out} \\sim r_{\\rm g}/r_0\\,$. Then, the total radiation efficiency of the disk region from $r = r_{\\rm H}$ to $r = r_\\infty$ would be $\\varepsilon_{\\rm total} = \\varepsilon_{\\rm in} + \\varepsilon_{\\rm out} \\sim r_{\\rm g}/r_0\\,$.  Finally, we remark that the issue of stability/instability has been ignored in our treatment. MHD instabilities are often important in the study of magnetized accretion disks (for reviews see \\cite{bal98,mes99,bla02}), they may also be important for our solutions. However, a detail treatment of the stability/instability of our solutions is beyond the scope of the current paper, so we would leave it for future work."
    },
    "0212/astro-ph0212029_arXiv.txt": {
        "abstract": "First results from a study into the abundance trends of oxygen in the Galactic thin and thick disks are presented. Oxygen abundances for 21 thick disk and 42 thin disk F and G dwarf stars based on very high resolution spectra ($R\\sim215\\,000$) and high signal-to-noise ($S/N>400$) of the faint forbidden oxygen line at 6300 {\\AA} have been determined. We find that $\\rm [O/Fe]$ for the thick disk stars show a turn-down, i.e. the ``knee'', at [Fe/H] between $-0.4$ and $-0.3$ dex indicating the onset of SNe type Ia. The thin disk stars on the other hand show a shallow decrease going from $\\rm [Fe/H] \\sim -0.7$ to the highest metallicities with no apparent ``knee'' present indicating a slower star formation history. ",
        "introduction": "The Galactic thin and thick disks are two distinct stellar populations in terms of age distributions and kinematics. The chemical trends in the two systems are also most likely different although recent works give conflicting results, see e.g. Chen et al.~(2000) and Fuhrmann~(1998). We show that the abundance trends for oxygen are different for the thin and thick disks. ",
        "conclusions": ""
    },
    "0212/astro-ph0212359_arXiv.txt": {
        "abstract": "We performed simulations of self-gravitating hydrodynamic turbulence to model the formation of filaments, clumps and cores in molecular clouds.  We find that when the mass on the initial computational grid is comparable to the Jeans mass, turbulent pressure is able to prevent gravitational collapse.  When the turbulence has damped away sufficiently, gravitational collapse can occur, and the resulting structure closely resembles the pre-singularity collapse of an isothermal sphere of \\cite*{penston69}.  If several Jeans masses are initially placed on the grid, turbulence may not be sufficient to prevent collapse before turbulence can be significantly damped.  In this case, the cores have density structures which are considerably shallower than expected for an isothermal gas, and resemble the solutions for a logatropic equation of state. ",
        "introduction": "There is abundant evidence that stars form in clumps within molecular clouds.  In order to understand the mechanism by which stars form, we therefore must understand the environment within a molecular cloud.  Observations show that molecular clouds are composed of a hierarchical network of filaments, clumps and cores (\\cite{mizuno95,johnstone99,gahm02}).  The most popular explanation for this substructure is that of turbulent fragmentation (\\cite{klessen00,ostriker01}).  In this scenario, turbulent fluctuations within the molecular cloud can compress the gas through shocks to high enough densities that the self-gravity of such regions becomes strong enough to induce collapse, resulting in a cloud which has been broken up into several condensing cores.  The cores, due to their origins as fragmenting shocks, are arranged along filamentary structures; the intersections of these filaments have many of these smaller cores attracted to their higher average density, resulting in clusters which resemble the observed clumps.  We performed simulations to test whether the structure of cores produced in this turbulent fragmentation scenario are consistent with observation. ",
        "conclusions": ""
    },
    "0212/astro-ph0212445_arXiv.txt": {
        "abstract": "Gamma-ray bursts (GRBs) are tremendous explosions visible across most of the Universe, certainly out to redshifts of $z=4.5$ and likely out to $z \\sim 10$.  Recently, GRBs have been found to have a roughly constant explosive energy as well as to have two luminosity indicators (the spectral lag time and the variability) that can be used to derive the burst's luminosity distance from the gamma-ray light curve alone.  There currently exists enough information to calibrate luminosity distances and independent redshifts for nine bursts.  From these, a GRB Hubble diagram can be constructed, where the observed shape of the curve provides a record of the expansion history of our Universe.  The current nine burst diagram is sparse, yet formal limits can be placed on the mass density of a flat Universe.  This first GRB Hubble diagram provides a proof of concept for a new technique in cosmology at very high redshifts.  With the launch of the SWIFT satellite in 2003, we should get $\\sim 120$ bursts to produce a Hubble diagram impervious to all effects of dust extinction and out to redshifts impossible to reach by any other method. ",
        "introduction": "As always in astronomy, the determination of distances is a crucial and difficult problem.  Until recently, the distance scale of GRBs was unknown by over 12 orders-of-magnitude.  In 1997, the discovery of optical and radio counterparts \\citep{cos97,van97,fra97} proved that at least the long-duration bursters were at cosmological distances with redshifts of $z \\sim 1$.  The measurement of GRB redshifts requires deep optical spectra, and to date only 24 redshifts are known for bursts with unknown selection effects.  If GRB distance indicators can be found that use only gamma-ray data, then we can measure the demographics and cosmology of large and well-understood samples of bursts.  Two such luminosity (and hence distance) indicators have recently been proposed.  The first \\citep{nmb00} relates the burst luminosity ($L$) with the spectral lag ($\\tau_{lag}$), which can be idealized as the time between peaks as recorded at high and low photon energies.  The second \\citep{frr00} relates the luminosity with the variability ($V$), which is a specific measure of the 'spikiness' of the burst light curve.  High luminosity bursts have short lags and spiky light curves, while low luminosity events have long lags and smooth light curves.  These two relations make GRBs into 'standard candles', in the same sense as for Cepheids and supernovae where an observed light curve property can yield the luminosity and then the distance.  The two luminosity indicators were originally proposed and calibrated with 6 or 7 bursts.  The addition of 3 or 2 further bursts (for a total of nine) have fallen on the original relations, hence adding confidence in their utility.  Further, if both the $L / \\tau_{lag} $ and $L/V$ relations are true, then there must be a particular $ \\tau_{lag} / V$ relation, and this prediction has been strongly confirmed with an independent sample of 112 BATSE bursts \\citep{sdb01}.  Finally, the very long lag bursts have been shown to be very low luminosity as demonstrated by their concentration to the local supergalactic plane \\citep{nor02}.  These luminosity indicators have been used to identify specific bursts \\citep{frr00} that are at redshifts of $z \\sim 10$ as well as to show that the star formation rate of the Universe is rising steadily \\citep{frr00,sdb01,lfr01} from $z \\sim 2$ to $z>6$.  This {\\it Letter} reports on the construction of a Hubble diagram (a plot of luminosity distance, $D_L$, versus redshift) as a means of measuring the expansion history of our Universe. ",
        "conclusions": ""
    },
    "0212/astro-ph0212115_arXiv.txt": {
        "abstract": "We analyze published reverberation mapping data for three Seyfert galaxies (NGC 3227, NGC~3516, and NGC~4593) to refine the mass estimate for the supermassive black hole in the center of each object. Treatment of the data in a manner more consistent with other large compilations of such masses allows us to more securely compare our results to wider samples of data, e.g., in the investigation of the M$_{bh}$-$\\sigma_{\\ast}$ relationship for active and quiescent galaxies. ",
        "introduction": "Reverberation mapping \\citep{bla82,pet93} is an important technique for probing the supermassive black holes (SMBHs) at the centers of active galactic nuclei (AGNs). Measuring the time lags, $\\tau$, between fluctuations in the highly variable AGN continuum and the variations of emission lines arising from the broad-line region (BLR) provides a responsivity-weighted radius for the BLR gas distribution around the SMBH. When emission-line time lags are combined with calculations of the gas velocity width, $\\sigma$, under the assumption of virialized gas motions, it becomes possible to estimate the SMBH mass.  Multiple lines of evidence suggest that to assume a virialized BLR is a reasonable approximation. In AGNs for which several emission lines have been reverberation-mapped (NGC~5548, 3C~390.3, NGC~7469, NGC~3783), an inverse relationship is found between $\\tau$ and $\\sigma$ that is consistent with expectations from virialization \\citep{pet99,pet00,onk02}. In addition, the close concordance between AGNs and quiescent galaxies in the $M_{bh}$-$\\sigma_{\\ast}$ plane under the virial assumption \\cite[relating black hole mass to the stellar velocity dispersion of the galactic spheroid;][]{fer01} would seem to imply that any systematic error in reverberation masses is small (a factor of $\\sim$3 or less). \\citet{kro01} estimated a similar level of expected systematic scatter assuming gravitational domination of the BLR.  Most of the statistical studies of reverberation-based black-hole masses use the homogeneous compilations of Wandel, Peterson, \\& Malkan (1999, hereafter WPM) and \\citet{kas00}. Missing from these compilations are AGNs observed by the ``Lovers of Active Galaxies (LAG)'' collaboration that monitored several AGNs in early 1990 \\citep{rob94}. Three of the LAG sources, NGC~3227 \\citep{sal94}, NGC~3516 \\citep{wan93}, and NGC~4593 \\citep{die94}, have well-determined emission-line lags and certainly meet the quality criteria for inclusion in these studies, and indeed were included in the compilation of Ho (1999; based on \\hb\\ only). These galaxies were omitted from the WPM compilation only because the data had not been analyzed in the same fashion as the other sources; specifically, (a) WPM use the model-independent Monte Carlo method of \\citet{pet98} to assess uncertainties in emission-line lags, and (b) WPM use the FWHM of the emission-line in the root-mean-square (rms) spectrum, rather than the mean spectrum, to characterize the BLR line-of-sight velocity width, $\\Delta V$, that is used to form the virial product ($\\Delta V$)$^{2} \\tau$. Because we have converted both our time lags and velocity widths to the AGN rest frame, our results supplement the compilation of \\citet{kas00}. (WPM did not correct for the time dilation of their $\\tau$ values.)  In this contribution, we reanalyze the LAG data on these three sources in the same fashion as \\citet{kas00}, with the intent of enlarging that homogeneous database. The total number of AGNs which have been reverberation-mapped (a few dozen) is small enough that additional objects will assist in on-going statistical investigations of (what are hoped to be) fundamental physical relationships. The main aspect on which we have worked is the careful construction the rms spectra over the spans of the respective observing campaigns. With the rms spectra for these AGNs and updated time lag values and uncertainties, we are able to derive masses for the SMBHs in these three galaxies.  In \\S\\ 2, we briefly describe the original observing campaigns and our data reduction. We explain the details of our time series and velocity width analyses in \\S\\ 3. Our estimates of the SMBH mass in each galaxy are presented in \\S\\ 4. In \\S\\ 5, we address the M$_{bh}$-$\\sigma_{\\ast}$ relationship, and we then summarize our conclusions (\\S\\ 6). ",
        "conclusions": "As many of the forefront studies of AGNs are examining correlations between SMBH masses and properties of the AGN host galaxies, a uniform analysis methodology is important for minimizing systematic offsets between datasets. Time lags, calculated from the published light curves, were used in conjunction with rms velocity widths for optical hydrogen emission lines to estimate the reverberation masses for NGC~3227, NGC~3516, and NGC~4593. These masses, listed in Table \\ref{tab4}, can now be added to the compilation of \\citet{kas00}, as they were calculated in the same manner. In addition, these masses, when combined with published data on the host galaxy spheroid velocity dispersions, strengthen the claim that reverberation masses are accurate measures of the true SMBH mass."
    },
    "0212/astro-ph0212522_arXiv.txt": {
        "abstract": "{ We collected a sample of 100 galaxies for which different observers have determined colour indices of globular cluster candidates. The sample includes representatives of galaxies of various morphological types and different luminosities. Colour indices (in most cases $(V-I)$, but also $(B-I)$ and $(C-T_1)$) were transformed into metallicities $\\rm [Fe/H]$ according to a relation by Kissler-Patig (1998). These data were analysed with the KMM software in order to estimate similarity of the distribution with uni- or bimodal Gaussian distribution. We found that 45 of 100 systems have bimodal metallicity distributions. Mean metallicity of the metal-poor component for these galaxies is ${\\rm \\langle [Fe/H]\\rangle} = -1.40 \\pm 0.02,$ of the metal-rich component ${\\rm \\langle [Fe/H]\\rangle} = -0.69 \\pm 0.03$. Dispersions of the distributions are 0.15 and 0.18, respectively. Distribution of unimodal metallicities is rather wide. These data will be analysed in a subsequent paper in order to find correlations with parameters of galaxies and galactic environment. ",
        "introduction": "In galaxy formation theories different physical processes may determine the formation and evolution of galaxies with different luminosities and different morphological types. Star formation is regulated by hierarchical clustering of dark matter halos, accretion of gas into dark matter potential well, reionisation of ISM by first stars, outflows of gas, galactic mergers, etc. On the other hand, despite varieties between galaxies it is important to find similarities between them. Globular cluster systems (GCSs) are known to exist in galaxies with rather different morphology (see e.g. Ashman \\& Zepf 1998). Therefore, it is interesting to compare the properties of GCSs in galaxies with different morphological types and different luminosities.  It has been firmly established that GCSs of many galaxies may have bimodal metallicity distribution (e.g. Ashman \\& Zepf 1998; Gebhart \\& Kissler-Patig 1999), and different scenarios have been proposed to explain bimodality. It is rather reasonable that mergers between galaxies may stimulate the formation of globular clusters (GCs) as well as dwarf galaxies (Schweizer 1987; Ashman \\& Zepf 1992; Weilbacher et al. 2000), but before accepting mergers as a dominating mechanism in creating bimodality, additional studies and arguments are needed. A multi-phase collapse process (Forbes et al. 1997) as a dominating mechanism needs additional theoretical calculations based on the general theory of the reionisation epoch in galactic formation. Accretion of metal-poor GCs from dwarf galaxies (C\\^ot\\'e et al. 1998) is also a rather plausible process, but before accepting it as a dominating process, additional studies are needed. Although the reasons of uni- or bimodality have not been fully understood, serious attempts have been made to explain bimodality in the case of some particular galaxies or particular group of galaxies (e.g. van den Bergh 1998; Harris et al. 2000; Forbes et al. 2001; Larsen et al. 2001a; Ashman \\& Zepf 2001).  In this paper, we compile a sample of colour distributions of GCS in galaxies. Colours were corrected from absorption in our Galaxy according to Schlegel et al. (1998) if not mentioned otherwise. Colours were analysed with the help of the KMM algorithm (Ashman, Bird \\& Zepf 1994). As initial data for the KMM we used individual cluster data and not the histograms of metallicity distributions. The KMM algorithm allows to determine the maxima in uni- or bimodal metallicity distributions. Correlations of derived metallicities with various properties of host galaxies and spatial distribution parameters of GCSs will be studied in a subsequent paper. ",
        "conclusions": ""
    },
    "0212/astro-ph0212464_arXiv.txt": {
        "abstract": " ",
        "introduction": "Our {\\GCR} provides, arguably, the best available laboratory in which to study the central activity of galaxies.  Recent measurements of the proper motion of infrared stars around the dynamical centre of the Galaxy coupled with the enhanced radio brightness of the region, strongly support the presence of a super massive black hole of mass (2--3)$\\times 10^6$ M$_\\odot$ ({\\eg} Genzel {\\etal} 2000) in Sgr~A$^*$.  A number of unusual phenomena are also observed in the {\\GCR}: {\\eg} magnetic radio filaments, high density, fast-moving molecular clouds, massive star clusters, and high-temperature plasma extending over scales in excess of a few hundred parsecs, etc.  A wide-angle {\\XMMN} survey of the {\\GCR} was carried out in the period 2000--2001 utilising 11 pointings in total, each with 10--25 ksec exposures (see Warwick 2002).  This {\\XMMN} {\\GC} survey along with a complementary {\\Chandra} survey (Wang {\\etal} 2002) has revealed the complex X-ray structure of the region.  Here we report some of the recent discoveries which have stemmed from the {\\XMM} {\\GC} Survey, focusing particularly on Sgr~A$^*$, Sgr~A East, the Radio Arc region, and an X-ray non-thermal filament designated XMM J174540$-$2904.5. ",
        "conclusions": ""
    },
    "0212/astro-ph0212378_arXiv.txt": {
        "abstract": "We present new results on the gravitational lensing shear and magnification power spectra obtained from numerical simulations of a flat cosmology with a cosmological constant. These results are of considerable interest since both the shear and the magnification are observables. We find that the power spectrum in the convergence behaves as expected, but the magnification develops a shot-noise spectrum due to the effects of discrete, massive clusters and symptomatic of moderate lensing beyond the weak-lensing regime. We find that this behaviour can be suppressed by ``clipping\" of the largest projected clusters. Our results are compared with predictions from a Halo Model-inspired functional fit for the non-linear evolution of the matter field and show excellent agreement. We also study the higher-order moments of the convergence field and find a new scaling relationship with redshift. In particular, the statistic $S_3$ is found to vary as $z_s^{-2.00\\pm 0.08}$ (where $z_s$ is the source redshift) for the cosmology studied, which makes corrections for different median redshifts in different observational surveys particularly simple to apply. ",
        "introduction": "Knowing the distribution and evolution of the large-scale structure in the universe, together with the cosmological parameters which describe it, are fundamental to obtaining a detailed understanding of the cosmology in which we live. Studies of the effects of weak gravitational lensing in the images of distant galaxies are extremely useful in providing this information. In particular, since the gravitational deflections of light arise from variations in the gravitational potential along the light path, the deflections result from the underlying distribution of mass, usually considered to be in the form of dark matter. The lensing signal therefore contains information about the clustering of mass along the line-of-sight, rather than the clustering inferred from galaxy surveys which trace the luminous matter.  Most obviously, weak lensing induces a correlated distortion of galaxy images. The magnitude of the correlations depends on the density parameter, $\\Omega_m$, and the value of the vacuum energy density parameter, $\\Omega_V$, for the universe, as these parameters reflect both the amount of mass and the rate of evolution of structure. Consequently, the correlations depend strongly on the redshifts of the lensed sources, as described by Jain \\& Seljak (1997) and Barber (2002). Recently a number of observational results have been reported for the so-called cosmic shear signal, which measures the variances in the shear on different angular scales. Bacon, Refregier \\& Ellis (2000), Kaiser, Wilson \\& Luppino (2000), Maoli et al. (2001), Van Waerbeke et al. (2000a, b), Wittman et al. (2000), Mellier et al. (2001), Rhodes, Refregier \\& Groth (2001), Van Waerbeke et al. (2001), Brown et al. (2002b), Bacon et al. (2002), Hoekstra, Yee \\& Gladders (2002), Hoekstra, Yee, Gladders, Barrientos, Hall \\& Infante (2002) and Jarvis et al. (2002) have all measured the cosmic shear and found good agreement with theoretical predictions.  In addition to shearing, weak gravitational lensing may cause a source at high redshift to become magnified or de-magnified as a result of the amount and distribution of matter contained within the beam. The degree of magnification is strongly related to the convergence, which represents a projection of the density contrast, $\\delta(x)$, at position $x$ along the line of sight, and which is proportional to $\\Omega_m$. Consequently, measurements of statistics for the convergence are able to provide cosmological constraints, and comparisons of the power spectrum of the convergence with theoretical predictions and numerical values serve to validate our theoretical and numerical models (see, e.g., Moessner \\& Jain, 1998, and Jain, 2002).  Of particular importance for interpreting weak lensing statistics is the fact that the scales of interest lie largely in the non-linear regime (see, e.g., Jain, Seljak \\& White, 2000). On these scales, the non-linear gravitational evolution introduces non-Gaussianity to the convergence distribution, and this signature becomes apparent in higher-order moments, such as the skewness. In addition, the magnitude of the skewness values is very sensitive to the cosmology, so that measurements of higher-order statistics in the convergence may be used as discriminators of cosmology.  In this work, we have obtained weak lensing statistics from cosmological $N$-body simulations using an algorithm described by Couchman, Barber \\& Thomas (1999) which computes the three-dimensional shear in the simulations. The code has been applied to cosmological simulations with $\\Omega_m = 0.3$ and $\\Omega_V = 0.7$; cosmologies of this type will be referred to as LCDM cosmologies.  To obtain the required statistics on different angular scales, the computed shear values have been combined (using the appropriate angular diameter distance factors and accounting for multiple deflections) along lines of sight arranged radially from the observer's position at redshift $z = 0$. Detailed results are presented for background sources at 14 different redshifts ($z_s = 0.1$ to 3.6) and angular scales from $1'$ to $32'$.  As a test of the accuracy of non-linear fits to the convergence power we compare the numerically generated convergence power spectra with our own theoretically predicted convergence spectra based on a Halo Model fit to numerical simulations (Smith et al., 2002). We also investigate the statistical properties of the magnification power spectrum and test predictions of the weak lensing regime. We also report on the expected redshift and scale dependence for higher-order statistics in the convergence.  A brief outline of this paper is as follows. In Section 2, we define the shear, reduced shear, convergence and magnification in weak gravitational lensing and outline how the magnification and convergence values are obtained in practice from observational data. In Section 3 we describe the relationships between the power spectra for the convergence, shear and magnification fluctuations, and how the power spectrum for the convergence relates to the matter power spectrum. We also describe our methods for computing the convergence power in the non-linear regime. Also in this Section, the higher-order moments of the non-linear convergence field are defined. The numerical procedure we use to generate the shear, convergence and magnification fields from the simulations are presented in Section 4, while in Section 5 we present our results for the numerical and theoretical comparison of the convergence power spectra, the power in the magnification fluctuations, and the higher-order moments, particularly the $S_3$ statistic. Finally we discuss the results and present our conclusions in Section 6. ",
        "conclusions": "\\subsection{Discussion}  We have shown that the convergence power spectrum values computed directly from the weak lensing simulations show remarkably good agreement with the non-linear predictions for the convergence power based on the Smith et al. (2002) Halo Model-inspired fitting formul\\ae. Whilst we have not compared the predictions for other cosmological models in this paper, the same numerical data, together with equivalent data for an open cosmology, have been compared with the predictions from Hierarchical models for the density field. Valageas, Barber \\& Munshi (2003) have shown excellent agreement for the full probability distribution functions for the shear and Barber, Munshi \\& Valageas (2003) have shown agreement with predictions for the convergence and higher-order moments. Consequently, these several agreements give us considerable confidence in our lensing procedure and the resulting weak lensing data, which we have here analysed in a number of ways.  Our results for the power in the magnification fluctuation suggest that the observed magnifications should be treated with some care as an estimate of the convergence power. This is especially true on small scales and at high redshift, since the magnification is very sensitive to locally high values of the convergence (arising from the tail of the skewed convergence distribution) and may vary significantly across small angular intervals. It further confirms that the weak lensing regime for the magnification should be treated with care, even for galaxy surveys with a mean redshift of 0.8. Worse, at a redshift of 1, where many surveys have been undertaken or planned, the magnification power may be extremely noisy. Consequently, evaluation of the convergence from magnification measurements is clearly likely to be biased by isolated high peaks in the convergence and shear fields.  The origin of this effect is the appearance of a few regions of very high magnification in the simulations. The spectrum is dominated by the strongest magnification events, although the signal shows moderate departures from the expected values even for medium magnifications. Since in an LCDM universe clusters form earlier and have time to relax, we expect large clusters to generate a large shear field. It is these largest structures with the highest shears, at the tails of the shear distribution, which dominate the magnification. These large shear-producing structures are rare, and so can be assumed to have a Poisson distribution. Hence the effect is to produce a shot-noise effect in the observed magnification power spectrum. Since the magnification is a non-linear function of the shear and convergence, this will not appear in the individual shear and convergence power spectra, but should be more apparent in their higher-order moments. Another consequence of this is that the cosmic variance, from realisation to realisation, should fluctuate quite wildly, which is the case here, and motivated the use of median statistics, rather than means which are more susceptible to this effect.  Since the magnification effect results from lensing from the largest structures and since the formation of structure evolves at different rates in different cosmologies, we would expect similar effects to be present in other cosmologies, but with differing degrees of severity. We have not studied the effect in other cosmologies but we might reasonably assume that the LCDM cosmology would display large effects because of the early development of structure.  We have shown that the effect can be damped by ``clipping\" the magnification field by removing the highest peaks. Removing the magnification values beyond 6-$\\sigma$ in the distribution removes most of the effect, showing that indeed it is the highest peaks causing this effect. One may wonder if this is a purely numerical effect, as in the simulations there is a high sampling of lines of sight, including the high-magnification regions around clusters. In reality we may not expect to see such extreme effects as often since galaxy surveys only Poisson sample the magnification field. In addition, if the normalisation, $\\sigma_8$, were lower than our chosen value, the number of clusters would be expected to be smaller, again leading to fewer examples of extreme magnifications.  However, even after clipping the magnification values beyond 6-$\\sigma$ there is still a residual systematic. Only after clipping the values beyond 1-$\\sigma$ in the distribution do we remove this effect significantly.  This all suggests that estimating the convergence power spectrum from magnification should be handled with care, unless evaluated at low redshifts. In addition, the estimation of higher-order statistics from the magnification will be equally affected by the shot-noise from individual, massive clusters which will also increase the effects of sampling variance on any large-scale measurement of the magnification field. Hence, while magnification remains a useful tool for probing the mass distributions in individual clusters, its value as a statistical probe over large scales may be complicated by these non-linear effects.  The results of Barber et al. (2000) for $S_3$ on large angular scales for different cosmological models, and the consistent results of Jain et al. (2000) clearly show the strong $\\Omega_m$ dependence of $S_3$. Consequently, provided the convergence field is obtained by reconstruction from the shear, the $S_3$ statistic may be used to good effect to discriminate different cosmologies. In determining $S_3$ from their simulations, Jain et al. (2000) quote values for sources at redshift 1 only. It is reassuring to note that the magnitudes of their $S_3$ values on the different angular scales, determined in quite a different way from ours (by reconstruction of the convergence from the shear data), are in good agreement at the smallest scales where non-linear effects are expected to be most pronounced. As the angular scale increases our values fall slightly more rapidly with scale.  Our expression for the angular scale and redshift dependence for $S_3$ clearly demonstrates the necessity of knowing the redshift distribution of the sources. This is the same conclusion reached by Barber (2002) who reported the angular scale and redshift dependences for the shear variance.  However, now that the redshift dependence of $S_3$ is established for our LCDM cosmolgy ($S_3 \\propto z_s^{-2.00}$), values determined from different surveys can be adjusted easily to a fixed median source redshift, enabling the combination and comparison of data from the different surveys. The only proviso here is that the redshift dependence of $S_3$ is likely to be cosmology dependent, since the magnitude of this statistic on a given angular scale and for a specific distribution of sources is known to be a discriminator of cosmologies. We have not yet investigated the redshift dependence of $S_3$ in other scenarios, although Barber et al. (2003) have reported on various aspects of $S_3$ from simulations and semi-analytical predictions in two different cosmologies. We would also hope to investigate the higher-order statistics for the convergence smoothed with a compensated filter in a later work, as this will allow more direct comparison with observed shear data.  The statistic $S_3\\sigma^2_\\kappa$ for a given angular scale is much less sensitive to the source redshift, as Bernardeau et al. (1997) also found for the Einstein-de Sitter cosmology. Figure~\\ref{s3} shows it to be approximately independent of redshift for high redshift sources and, using the results for the redshift dependence of the shear variance from Barber (2002) together with the redshift dependence of $S_3$ reported here, the statistic is shown to be only a very slowly increasing function of $z_s$ for $z_s < 1$. Consequently, $S_3\\sigma^2_\\kappa$ may prove to be very useful for surveys in which the redshift distribution of the sources is uncertain.  \\subsection{Conclusions}  For the convergence power spectrum, $\\ell(\\ell+1)C^{\\kappa\\kappa}_\\ell/(2\\pi)$, computed numerically from the lensing statistics in our simulations, we find excellent agreement with the values computed using a halo model. We can therefore have considerable confidence in our weak lensing statistics when applied to the shear, convergence, magnification and higher-order statistics. In addition, with this consistency, theoretical determinations of lensing statistics on a wider range of scales than can be achieved with numerical simulations may be used with increasing confidence.  Our results for one quarter the power in the magnification fluctuations, $\\ell(\\ell+1)C^{\\delta \\! \\mu \\delta \\!\\mu}_\\ell/(8 \\pi)$, show that the magnification is susceptible to the effects of discrete massive clusters and large variations across small angular intervals. These effects occur specifically beyond the weak lensing regime, and become apparent in the magnification even at low redshift and on small angular scales. Consequently, determination of the convergence field from magnification data should be treated with special attention.  Our simple mathematical description for $S_3(\\theta, z_s)$, showing it to be closely proportional to $z_s^{-2}$ for the LCDM cosmolgy, makes it particularly simple to compare and combine the results from different surveys in which the median redshifts of the galaxies may be different.  Finally, we found directly, and through combination with the mathematical expressions for the shear variance and $S_3$, that the combined statistic $S_3\\sigma^2_\\kappa$ is only a very slowly increasing function of redshift for low redshift sources and approximately independent of redshift at high redshift, as has been found in the Einstein-de Sitter cosmology. This statistic, therefore, may be useful in comparing and combining the data from different surveys with different galaxy redshift distributions."
    },
    "0212/hep-ex0212034_arXiv.txt": {
        "abstract": "The detection of low energy neutrinos ($<$ few tens of MeV) via coherent nuclear scattering remains a holy grail of sorts in neutrino physics. This uncontroversial mode of interaction is expected to profit from a sizeable increase in cross section proportional to neutron number squared in the target nucleus, an advantageous feature in view of the small probability of interaction via all other channels in this energy region. A coherent neutrino detector would open the door to many new applications, ranging from the study of fundamental neutrino properties to true \"neutrino technology\". Unfortunately, present-day radiation detectors of sufficiently large mass ($>$ 1 kg) are not sensitive to sub-keV nuclear recoils like those expected from this channel. The advent of Micropattern Gas Detectors (MPGDs), new technologies originally intended for use in High Energy Physics, may soon put an end to this impasse. We present first tests of MPGDs fabricated with radioclean materials and discuss the approach to assessing their sensitivity to these faint signals. Applications are reviewed, in particular their use as a safeguard against illegitimate operation of nuclear reactors. A first industrial mass production of Gas Electron Multipliers (GEMs) is succinctly described. ",
        "introduction": "\\PARstart{A} new family of radiation detector designs, generally referred to as Micropattern Gas Detectors (MPGDs) \\cite{reviews} has emerged during the last fifteen years in response to the demanding needs (fast counting rate, radiation resistance, high spatial resolution) of next-generation High Energy Physics experiments. While the specific design varies, their common principle is a sizeable voltage drop across microstructures immersed in a suitable gas mixture: electrons originating from particle ionization in a conversion volume are multiplied in the microstructures, where amplification gains of up to $10^{7}$ are obtained. A popular example of a MPGD is the MICROMEGAS design (MICROMEsh GAseous Structure), a concept recently put forward by Y. Giomataris, G. Charpak and collaborators \\cite{mms1}. This two-stage parallel-plate avalanche chamber consists of a 100 $\\mu$m narrow amplification gap and a large conversion region -TPC volumes are possible-, separated by a gauze-like electroformed conducting micromesh. Electrons released by ionizing particles in the gas-filled conversion region are drifted towards the amplification gap where they multiply in an avalanche process. Detectable signals are then induced on anode elements. A second example of MPGDs are Gas Electron Multipliers (GEMs)\\cite{GEMS}, developed at CERN by F. Sauli and collaborators: small holes (diameter$\\sim 80 \\mu m$) are photolithographically etched on a $\\sim 50 \\mu m$-thick Kapton film copper-clad on both sides and a voltage difference of $\\sim 400$ V is generated across the GEM. The high density of electric field lines within the perforations induces the sought avalanche. An advantage of GEMs is the possibility of building multi-stage amplification layers \\cite{ian}, allowing for very large gains. The high-efficiency detection of single electrons at gas pressures of up to 20 atm has been achieved in a variety of MPGDs \\cite{mms5}. The effective energy threshold in these devices is the ionization energy of the gas mixture, i.e., a few tens of eV. \\begin{figure}[tbp] \\epsfxsize = \\hsize \\epsfbox{fig1.eps} \\caption{Detectable signal in different gases from neutral-current nuclear scattering of reactor antineutrinos (10$^{13}  \\bar{\\nu}$ cm$^{-2}$ s$^{-1}$), obtained by folding of the differential cross section in [8] with the reactor spectrum in [24] and applying quenching factors derived from SRIM [25]. The tradeoff between endpoint energy and rate with increasing atomic mass is evident. {\\it Table}: total coherent recoil rate in different gases under the same conditions.} \\end{figure} \\begin{figure}[tbp] \\epsfxsize = \\hsize \\epsfbox{fig2.eps} \\caption{Distribution of the number of ionized electrons produced by coherent nuclear scattering in a gaseous detector exposed to a typical reactor antineutrino flux. This estimate includes the reactor emission spectrum, differential cross section, a quenching factor derived from Linhard's theory and the mean ionization energy of the gas mixtures. {\\it Top:} for pure noble gases, {\\it Bottom:} after addition of a small fraction of TMAE vapor which may in principle reduce the ionization threshold to $\\sim$6 eV. The effect of gas additives such as TMAE or TEA on energy threshold is to be investigated as part of this work.} \\end{figure}  The possibility of exploiting some of the features specific to MPGDs in a new realm, that of searches for rare-events in neutrino and astroparticle physics has been recently examined \\cite{meyannis}. The properties of these devices (background rejection capabilities, demonstrated ability for single-electron detection, versatility and simplicity) suggest a means to tackle a long-standing experimental challenge, the measurement of {\\it coherent} neutral-current neutrino-nucleus scattering. An uncontroversial process in the Standard Model, the scattering off nuclei of low-energy neutrinos ($< $ few tens of MeV, e.g., reactor $\\bar{\\nu} s$) via the neutral current \\cite{freedman} remains undetected. The long neutrino wavelength probes the entire nucleus, giving rise to a large coherent enhancement in the cross section, roughly proportional to neutron number squared \\cite{drukier}. Using this mode of interaction, it would be possible to speak of {\\it portable} neutrino detectors: in some experimental conditions the expected rates can be as high as several hundred recoils/kg/day, by no means a ``rare-event'' situation. However, the recoil energy transferred to the target is a few keV at most even for the lightest nuclei, with only a few percent going into ionization (\\frenchspacing{Figs. 1,2}).  The interest in observing this process is not merely academic: a neutral-current detector responds the same way to all known neutrino types. Therefore, the observation of neutrino oscillations in such a device would be {\\it direct} evidence for a fourth sterile neutrino. These can be invoked if all recently observed neutrino anomalies are accepted at face value \\cite{ellis} and may play an important role as Dark Matter \\cite{dolgov}. Separately, the cross section for this process is critically dependent on neutrino magnetic moment. Agreement with the Standard Model prediction would {\\it per se} largely improve on the present experimental sensitivity to $\\mu_{\\nu}$\\cite{dodd}. In addition to this, a measurement of the cross section would constitute a sensitive probe of the weak nuclear charge, testing radiative corrections due to new physics above the weak scale with a sensitivity comparable to atomic parity violation and accelerator experiments. Statistically speaking, this can be accomplished in a nuclear reactor already with a modest detector mass and a short exposure \\cite{larry}. Finally, this coherent mechanism plays a most important role in neutrino dynamics in supernovae and neutron stars \\cite{freedman}, adding to the attraction of a laboratory measurement of this cross section. In particular, a measurement of the {\\it total} (flavor-independent) neutrino flux from a nearby supernova using a large enough coherent detector would be of capital importance to help clarify the exact oscillation pattern followed by the neutrinos in their way to the Earth \\cite{beacom}.  \\begin{figure}[tbp] \\epsfxsize = \\hsize \\epsfbox{fig3.eps} \\caption{Power and approximate antineutrino flux distribution of worldwide nuclear reactors, extracted from the databases in [26]. Proposals to monitor illicit reactor activity using conventional neutrino detectors (large liquid scintillator tanks) seem insufficient for this purpose: their sensitivity under realistic conditions would be adequate only for reactor powers larger than $\\sim 3$ GWt and require the construction of underground infrastructure ($\\sim$ 6 m.w.e.) close to reactor cores [16]. Detectors based on coherent scattering may be able to improve this situation in the near future.} \\end{figure} \\begin{figure}[tbp] \\epsfxsize = \\hsize \\epsfbox{fig4.eps} \\caption{{\\it Top}: A first mass-production of GEMs using 3M's Microflex adhesiveless reel-to-reel process. The roll in the figure contains 35 panels of 33 GEM elements each. Any GEM pattern up to 12''x12'' can be produced. Perforations to facilitate detachment are visible around each element. {\\it Bottom right}: SEM photograph of one of these GEMs (hole diameter 80 $\\mu$m, pitch 140 $\\mu$m). {\\it Bottom left}: LEMs produced at EFI using automated micromachining on low background laminates (OFHC copper plated directly onto virgin Teflon).} \\end{figure} \\begin{figure}[tbp] \\epsfxsize = \\hsize \\epsfbox{fig5.eps} \\caption{Preliminary tests of gas gain in a single LEM  (Ar:DME(9:1) at 1 atm) using a $^{55}$Fe uncollimated source. The drift cathode was grounded and the LEM held at a positive potential (drift distance = 0.5 cm). A multi-layer LEM structure should provide enough amplification to detect single electrons at moderate gas overpressures. We plan to investigate the dependence of the proportional-scintillation light yield on operating pressure as a possible mechanism to increase gas density while maintaining single-electron sensitivity.} \\end{figure}  Until now, no existing device had met the mass and energy threshold requirements involved in this measurement, even though unrealized cryogenic proposals abound \\cite{blas}. A considerable fraction of the neutrino signal in a reactor experiment is nevertheless expected above MPGD energy thresholds (Fig. 2). Structurally simple MPGD-based coherent neutrino detectors would open the door to more mundane but no less important applications than those listed above (``neutrino technology''? \\cite{leocastle}): for instance, nonintrusive monitoring of nuclear reactors against illegitimate uses (e.g., fuel rod diversion, unauthorized production of weapon-grade material) with a compact device, potentially improving on existing proposals that rely on standard neutrino detectors and processes \\cite{bernstein} (Fig. 3). ",
        "conclusions": ""
    },
    "0212/astro-ph0212134_arXiv.txt": {
        "abstract": "We calculate the evolution of a low-mass ($M \\le 10^5 M_{\\odot}$) spherically symmetric density perturbation in the $\\Omega_b h^2=0.02$, $\\Omega_M =0.35$, $\\Omega_{\\Lambda}=0.65$, $h=0.72$ Universe. The results are compared with the ones that assume no cosmological constant and the flat, dark matter dominated Universe. We include thermal processes and non-equilibrium chemical evolution of the collapsing gas. We find that direct formation of bound objects with such masses by $z=8$ is unlikely so in fact they may form only through fragmentation of greater objects. This is in stark contrast to the $\\Omega=1$ pure CDM cosmology, where low-mass objects form abundantly at redshits $z>10$. ",
        "introduction": "Present observations of the CMB anisotropies by the BOOMERANG \\citep{Net02}, MA\\-XI\\-MA-1 \\citep{Lee01}, CBI \\citep{Pad01} and DASI \\citep{Pry02} experiments combined with the data from the Supernova project (\\citet{Rie98}; \\citet{Per99}) suggest that although the geometry of the Universe is flat, its matter content is dominated by the vacuum energy or so-called quintessence. The background cosmology has a very strong influence on the structure formation. In this paper we address the formation of first bound objects in the flat $\\Lambda$CDM cosmology.  The structure formation on both large and small scales has been studied very carefully in the dark matter dominated cosmology, in particular in the flat pure CDM model. This is not so for the vacuum-energy-dominated models, which are of prime relevance according to the above observations. Perhaps one of the most interesting objects to study are the first bound objects in the Universe -- the first quasars and the Population III stars that were the first sources of photons after the light of CMB faded away, and ended the `Dark Ages' of the Universe. Some of these calculations were described in a review by \\citet{Bar01}. The recent identification of quasars at $z\\sim6$ in the Sloan Digital Sky Survey \\citep{Fan01} strongly motivates such an investigation.  Our code is similar to the one by \\citet{Tho95} and by \\citet*{Hai96}. We consider the evolution of a single spherically symmetric density perturbation  in the early Universe starting soon after recombination until it finally forms a bound stationary object. Our aim here is to extend their analysis, which was restricted to the $\\Omega=1$ CDM cosmology, to the $\\Lambda$-dominated Universe and compare the results with  those for the flat, dark matter dominated Universe but with more up-to-date cosmological parameters ($h=0.72$, $\\Omega_b h^2=0.02$).  We start tracing the initial expansion of the perturbation at high redshift when its density contrast is still in the linear regime. Then we follow decoupling of the perturbation  from the Hubble flow and its subsequent collapse and formation of a virialized cloud. We include the gas dynamics and various cooling and heating processes operating in the expanding and collapsing cloud. The chemical evolution of the collapsing primordial gas cloud is also accounted for.  After the initial collapse, a virialized gas cloud is formed. The kinetic energy of the infalling gas is dissipated through shocks and the cloud becomes pressure-supported. We study the further evolution of the cloud which is determined by its ability to cool sufficiently fast. The most important cooling mechanism for low-mass clouds is the radiation of excited H$_2$ molecules. The presence of a small amount of the molecular hydrogen H$_2$ is crucial for triggering the final collapse of such clouds which could form the first luminous object in the Universe. ",
        "conclusions": "For the dark energy dominated Universe we need very large overdensities, usually much exceeding the 3$\\sigma$ limit, for low-mass ($M_0 \\leq 10^5 M_{\\odot}$) objects in order to at least partially collapse before $z=8$. For the flat, dark matter dominated Universe the collapse is much faster and direct formation of a low-mass object is quite likely, even as soon as for $z\\sim 20 - 30$. It is a result of much lower normalization in $\\Lambda$CDM (Table 1) and lower virial temperatures for $\\Lambda$CDM clouds of the same mass and initial overdensity. A secondary reason may be that in the flat CDM model any overdensity will stop its evolution and possibly cool and collapse while for $\\Omega_M < 1$ there is some threshold overdensity that depends on the model and epoch. The chemical evolution is pretty much insensitive to cosmology, at least before the final collapse. In the $\\Lambda$-dominated Universe perhaps low-mass objects may form through fragmentation of greater objects. Direct formation of such objects before $z\\sim 10$ seems  unlikely."
    },
    "0212/astro-ph0212072_arXiv.txt": {
        "abstract": "We combine new VLA D array HI data of NGC~2613 with previous high resolution data to show new disk-halo features in this galaxy. The global HI distribution is modeled in detail using a technique which can disentangle the effects of inclination from scale height and can also solve for the average volume density distribution in and perpendicular to the disk.  The model shows that the galaxy's inclination is on the low end of the range given by Chaves \\& Irwin (2001) and that the HI disk is thin ($z_e$ = 188 pc), showing no evidence for halo.  Numerous discrete disk-halo features are observed, however, achieving $z$ heights up to 28 kpc from mid-plane.  One prominent feature in particular, of mass, $8\\,\\times\\,10^7$ M$_\\odot$ and height, 22 kpc, is seen on the advancing side of the galaxy at a projected galactocentric radius of 15.5 kpc.  If this feature achieves such high latitudes because of events in the disk alone, then input energies of order $\\sim$ 10$^{56}$ ergs are required. We have instead investigated the feasibility of such a large feature being produced via buoyancy (with drag) within a hot, pre-existing X-ray corona. Reasonable plume densities, temperatures, stall height ($\\sim$ 11 kpc), outflow velocities and ages can indeed be achieved in this way. The advantage of this scenario is that the input energy need only be sufficient to produce blow-out, a condition which requires a reduction of three orders of magnitude in energy.  If this is correct, there should be an observable X-ray halo around NGC~2613. ",
        "introduction": "\\label{introduction}  High latitude HI in disk galaxies is increasingly being recognized as an important tracer of disk activity and dynamics. The gas scale height is a measure of time integrated activity in the disk and discrete features trace the level of local activity. The disk-halo interface represents a severe transition between two very different environments:  the high density, low velocity dispersion disk and the low density, high velocity dispersion halo.  Globally, it may provide information on the shape of the dark matter halo (e.g. Olling \\& Merrifield 2000). In a few recent cases, lagging HI halos have also been detected, such that the high latitude gas is rotating $\\sim$ 20 - 50 km s$^{-1}$ more slowly than the underlying disk (Schaap, Sancisi, \\& Swaters 2000, Fraternali et al. 2001, Lee et al. 2001, Rand 2000 and T{\\\"u}llmann et al. 2000). Explanations for such lags have focussed mainly on galactic fountains (cf. Bregman 1980, Swaters et al. 1997, Schaap et al. 2000, Collins et al. 2002).  Edge-on galaxies are particularly useful in the study of the disk-halo region and, in this paper, we especially wish to investigate how discrete disk-halo features can be used as probes of the environment into which they are emerging.  The edge-on galaxy, NGC2613 (D = 25.9 Mpc) is a good candidate for such a study.  In Paper I (Chaves \\& Irwin 2001), we presented Very Large Array (VLA) C Array observations of this galaxy.  The resulting C array zeroth moment map over an optical image is shown in Fig.~\\ref{Carray}, illustrating the presence of six symmetrically placed vertical extensions in this galaxy. In this paper, we present new VLA D array HI data of NGC~2613 and examine them separately as well as in combination with the previously obtained C array data.  The observations are presented in Sect.~\\ref{observations}, the results, including a description of extremely large high latitude HI features and a model of the global HI distribution are given in Sect.~\\ref{results}, in Sect.~\\ref{discussion}, we present the discussion, including a feasibility study of a buoyancy model, and a summary is given in Sect.~\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  New VLA D array data of NGC~2613 have been combined with previous higher resolution observations (Chaves \\& Irwin 2001) to show a more extensive HI distribution than previously observed. The galaxy is now seen to have a tidal tail on its eastern side due to an interaction with its companion, ESO~495-G017, to the north-west.  The three-dimensional HI distribution in NGC~2613 has been modeled following Irwin \\& Seaquist (1991) and Irwin (1994), a method which allows the volume density distribution to be determined as well as the scale height and inclination to be disentangled.  We find that the inclination of the galaxy (79$^\\circ$) is on the low end of the range given in Chaves \\& Irwin (2001) and the model now shows that there is no HI halo in NGC~2613.  Rather, the global HI distribution is well fit by a thin disk of exponential scale height, $\\it z_e$ = 188 pc. The use of such a model is very important in drawing conclusions about the presence or absence of a global halo in a galaxy of this inclination. Previous reports of a lagging halo from PV slices alone can largely be attributed to projection against the background disk.  While there is no significant global HI halo in NGC~2613, there are  more discrete  disk-halo HI features than previously detected and these HI features achieve extremely high latitudes.  Even though a tidal interaction is occurring, we suggest that most of the discrete kpc-scale features have been produced internally rather than from impacting clouds, although we do not rule out the possiblity of the companion galaxy having some indirect effect (e.g. triggering instabilities). The presence of many discrete features may be related to the fact that the global HI disk is thin, favouring blow-out. The observed  {\\it z} heights are quite remarkable (e.g. up to 28 kpc). The -307 km s$^{-1}$ feature, in particular, below the plane  on the advancing side, reaching 22 kpc in $\\it z$ height and of total mass, (8 $\\pm$ 2) $\\times$ 10$^7$ M$_\\odot$, is very obvious and well-defined in PV space. Its center is likely close to its projected radius of 15.5 kpc and it extends over a large ($\\pm$ 7 kpc radius) projected galactocentric radius. If this feature has achieved its $\\it z$ height as a result of internal processes, then extremely large energies are required, $\\sim$ 10$^{56}$ ergs.   Given the very high input energies required for the -307 km s$^{-1}$ feature, its resemblance to smaller buoyant features (cf. the Galactic Mushroom, English et al. 2000), and the fact that X-ray halos are being found around an increasing number of star forming spiral galaxies, we have carried out a feasibility study as to whether this feature can be interpreted as an adiabatic buoyant plume. The observed HI would be carried out by a hot, low density outflowing gas and, after having achieved blowout, would rise through a hot pre-existing X-ray corona. A reasonable example (Model 5), has a mean plume temperature and density of  $1\\,\\times\\,10^7$ K and 5.5 $\\times$ 10$^{-3}$ cm$^{-3}$ rising into a hot isothermal corona of temperature and mean density, $5\\,\\times\\,10^6$ K and 0.035 cm$^{-3}$, respectively. These conditions produce a stall height of 11 kpc which is where the observed plume widens in velocity and position space. The coronal density at the stall height is 2.8 $\\times$ 10$^{-3}$ cm$^{-3}$. The maximum outflow velocity in this model is 290 km s$^{-1}$ and it reaches the stall height in 5.4$\\,\\times\\,$10$^7$ yrs.  This model shows that, even with buoyancy alone (and there may be additional sources of pressure such as magnetic fields), HI can reach these extreme {\\it z} heights.  The advantage of such a model is that the energy requirements from the initial event are drastically reduced to being only what is required for blowout, a reduction of several orders of magnitude.  The energy is largely extracted from the gravitational potential of the galaxy rather than the initial event within the disk.  The behaviour of the plume should sample the parameters of the X-ray corona. The predicted  X-ray luminosity suggests that the corona should be observable at heights below the stall height although we expect that the distribution of X-ray emission may not be as smooth as assumed by the model."
    },
    "0212/astro-ph0212558_arXiv.txt": {
        "abstract": "The spherically symmetric stationary transonic (Bondi) flow is considered a classic example of an accretion flow. This flow, however, is along a separatrix, which is usually not physically realizable. We demonstrate, using a pedagogical example, that it is the dynamics which selects the transonic flow. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/astro-ph0212250_arXiv.txt": {
        "abstract": "Characterization of extrasolar planetary systems requires a radial velocity (RV) survey for planets around hundreds of thousands of nearby stars of all spectral types over next ten years. This survey will be  extremely difficult to conduct using current high resolution echelle spectrometer due to its single object observing mode and low instrument throughput. Here we propose to use a high throughput multi-object dispersed fixed-delay interferometer  for the survey. This instrument, a combination of  a fixed-delay interferometer with a moderate resolution spectrometer, is completely different from current echelle spectrometers. Doppler RV  is measured through monitoring interference fringe  shifts of stellar absorption lines over a broad band. Coupling this  multi-object instrument with a wide field telescope (a few degree, such as Sloan and WIYN) and UV, visible and near-IR detectors will allow to simultaneously obtain hundreds of stellar fringing spectra for searching for planets. The RV survey speed can be increased by more than 2 orders of magnitude over that for the echelles.   A prototype dispersed fixed-delay interferometer has been observed at the Hobby-Eberly 9m and Palomar 5m  telescopes in 2001 and demonstrated photo noise limited Doppler precision with Aldebaran. Our recent observations at the KPNO 2.1m telescope in 2002 demonstrate a short term Doppler precision of $\\sim$ 3 m/s with  $\\eta$ Cas (V = 3.5), a RV stable star and also obtained a RV curve for 51 Peg. (V = 5.5), confirming previous planet detection with an independent RV technique. The total measured detection efficiency including the sky, telescope and fiber transmission losses, the instrument and iodine transmission losses and detector quantum efficiency is 3.4{\\%} under 1.5 arcsec seeing conditions, which is comparable to all of the echelle spectrometers for planet detection. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/astro-ph0212299_arXiv.txt": {
        "abstract": "CCD photometry on the intermediate-band $uvbyCa$H$\\beta$ system is presented for the old open cluster, NGC 6253. Despite a high level of field star contamination due to its location toward the galactic center, combination of the data from the multiple color indices with the core cluster sample derived from radial star counts leads to the identification of a set of highly probable, single cluster members. Photometric analysis of a select sample of 71 turnoff stars produces a reddening value of $E(b-y) = 0.190 \\pm 0.002$ (s.e.m.) or $E(B-V)$ = 0.260 $\\pm 0.003$ (s.e.m.) from 71 stars.  The metallicity indices, $\\delta$$m_1$ and $\\delta$$hk$, both identify this cluster as the most metal-rich object studied on either system to date. Simple extrapolation of the available metallicity calibrations leads to [Fe/H] ranging from +0.7 to +0.9. Metal-rich isochrones with overshoot imply an age between 2.5 and 3.5 Gyr, with an apparent distance modulus between $(m-M)=$ 11.6 and 12.2, depending upon the isochrones used.  The improvement in the fit using $\\alpha$-enhanced isochrones may indicate that the cluster [Fe/H] is closer to +0.4, but the photometric indices are distorted by an elemental distribution other than a scaled solar. The galactocentric position of the cluster, in conjunction with data for other clusters and Cepheids, is consistent with the inner disk reaching and maintaining a metallicity well above solar since the early history of the disk, unlike the solar neighborhood. ",
        "introduction": "The shared properties of distance, age, and composition for stars within a cluster supply justification for their study within the context of stellar and galactic evolution. An example of how a well-defined and homogeneously-analyzed cluster sample can shed light on galactic evolution is provided by \\citet{TAT97}, wherein analysis of over 70 open clusters demonstrates the existence of structure within the galactic abundance gradient, structure that has been corroborated most recently through the use of Cepheids by \\citet{AN02}. Despite the unusual combination of high precision and sample size for the collected data, critical regions of the galactic disk remain inadequately sampled. Moreover, only a modest fraction of the open clusters within observational reach have their fundamental parameters determined to a degree of reliability that can satisfy the needs of Galactic studies tied to the global characteristics of the population.  Though the number of broad-band photometric surveys of individual clusters with color-magnitude diagrams (CMD) of sufficient precision has increased significantly in the last five years with the expanded availability of larger-format CCD's, the interpretation of these data has been hampered by the absence of a comparable progress in determinations of the critically important properties of reddening and metallicity, linchpins to age and distance. Attempts to constrain the parameters using CMD morphology and comparison to theoretically synthesized CMD's have had some success \\citep{TO91,BO90,GO96,MA97}, though the answers drawn from different models do not always concur and the uncertainties in the parameters remain larger than desirable for many projects.  To improve this situation for individual clusters and to expand and test the sample delineating the galactic structure outlined by \\citet{TAT97}, a program of CCD photometry of open clusters on the extended Str\\\"{o}mgren intermediate-band ($uvbyCa$) and H$\\beta$ systems has been undertaken. Examples of the initial work illustrating the value and precision attainable with careful analysis have been given for the metal-rich open cluster, IC 4651 \\citep[hereinafter referred to as Paper I]{AT00a} and the globular cluster, NGC 6397 \\citep[hereinafter referred to as Paper II]{AT00b}. These clusters were selected to sample the extreme range of cluster parameters with objects that had been reasonably well-studied by other photometric and spectroscopic means. Our first program cluster with no intermediate-band photometric work or reliable high-resolution spectroscopic abundances is the distant open cluster, NGC 6253.  NGC 6253 has been the focus of three broad-band photometric surveys to date, \\citet{BR97}, \\citet{PI98}, and \\citet{SA01} (hereinafter referred to as BR, PI, and SA, respectively). It is of interest because, as carefully demonstrated in BR, the cluster is located toward the galactic center ($l^{II} = 336\\arcdeg$, $b^{II} = -6\\arcdeg$), significantly older than the Hyades, and super-solar in metallicity, while exhibiting a richly-populated turnoff and giant branch. Despite the concurrence of opinion regarding the broad features of the cluster, examination of previous work reveals a wide range for the key parameters of interest. Previously published work produces plausible reddening estimates ranging from $E(B-V)$ = 0.20 to 0.32, metallicity from solar abundance to almost triple the solar value, and an age from 2 to 5 Gyr; we will discuss the specifics of each property and its past derivation in Sec. 4. Since traditional linear abundance gradients imply a change in the mean [Fe/H] of less than 0.3 dex over a 4 kpc galactocentric range, definitive evaluation of any structure within 2 kpc of the sun requires [Fe/H] data of significantly higher accuracy than is currently available for NGC 6253, for a large number of clusters. Moreover, a definitive metallicity is important in defining the value of the cluster as a template for younger, metal-rich, extragalactic populations and as a potential source of stars with planetary systems.  The underlying basis for the use of intermediate-band photometry is simple.  The brighter observational limit imposed by the narrower bandpasses is, in part, compensated by the higher precision with which one may disentangle the stellar parameters. DDO photometry \\citep{TAT97} recommends itself because it focuses on the brightest stars within a cluster, the red giants, and the calibration of the system is exquisitely tied to the spectroscopic abundances of the extensive sample of field giants near the sun \\citep{TAT96}. The weaknesses of the system are the small sample of giants often present in many poorly populated clusters observed within a field rich with cluster non-members and the possibility that post-main-sequence evolution may alter the atmospheric abundances of some stars.  The extended Str\\\"{o}mgren and H$\\beta$ systems exploit the rich population of main sequence stars and have been well-calibrated for spectral types ranging from B through G \\citep{CR75,CR78,CR79,OLS88,SN89}, the temperature range typified by the turnoff in most open clusters. No open cluster has been identified that exhibits a range in [Fe/H] or a difference in [Fe/H] between the turnoff stars and the red giants when the abundances are based primarily upon the Fe lines. The larger population of main sequence stars can lead to easier identification of a probable sample of cluster members while reducing the standard error of the mean for the final cluster parameters. The multiple indices also enhance the elimination of probable non-members from the final sample. For an example of how successful this approach can be for a typical open cluster, the reader is referred to Paper I.  Section 2 contains the details of the 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 theoretical isochrones. Sec. 6 summarizes our conclusions in the context of our current understanding of the evolution of the galactic abundance gradient. Preliminary results for NGC 6253 presented in \\citet{TAD} are superseded by this discussion. ",
        "conclusions": "The open cluster, NGC 6253, has unusual significance in two ways. It populates a rather rare class of open clusters, old and significantly closer to the center of the Galaxy than the sun. The statistical studies of \\citet{JP94} and \\citet{HUF} have demonstrated that, observational difficulties aside, the number of open clusters observed within 7 kpc of the galactic center is so small that the probability for survival beyond 1 Gyr is low; NGC 6253 falls just outside this zone. Second, only the old open cluster NGC 6791 has been identified as having a metallicity significantly higher than the Hyades \\citep{CHA}, and even this claim has been questioned \\citep{TAY}. Past attempts to quantify the key parameters for NGC 6253 have depended primarily on consistency fits to multicolor, broad-band relations and synthetic CMD's tied to theoretical isochrones. Though there is a consensus linked to the CMD morphology that the cluster must be 2 to 4 Gyr old and more metal-rich than the sun, narrowing the range has been difficult. Though the source of the disagreement has originated mostly with the different predictions from the various isochrones, a problem repeated in Sec. 5, the limitations imposed by the differences among the photometric surveys must be accounted for. As noted in the Appendix, the absolute answer one derives for the parameters would depend upon which zero-point and which color slope one chooses as the $correct$ system for defining the broad-band CMD and color-color plots.  We reiterate one of the chief advantages of the intermediate-band approach because it allows to fix the reddening through H$\\beta$ with only a weak dependence on the cluster metallicity. Removing this as a free parameter in any match to theoretical isochrones significantly restricts the allowed combinations of metallicity, distance, and age. As has now been demonstrated using IC 4651 (Paper I), NGC 6397 (Paper II), and NGC 6253, it is possible to derive the distance with an internal error of only $\\pm 0.1$ mag, with internal errors of the mean for $E(b-y)$ and [Fe/H] of $\\pm 0.002$ and $\\pm 0.02$, even though we cannot guarantee that all non-members and binaries have been eliminated.  Given this precision, how does one explain the extraordinary metallicities implied by both the $m_1$ and $hk$ indices? The flip side to the internal precision that is achievable is that the primary source of uncertainty in deriving photometric-based parameters from a large statistical sample falls upon the transformations of the instrumental indices to the standard system. The challenge of doing this well, even for an established broad-band system like UBV, is obvious from the comparisons of the data on NGC 6253 as given in the Appendix. However, as noted in Sec. 4, to simultaneously reduce the cluster metallicity by claiming an error in the photometry requires either a conspiracy between $m_1$ and $hk$ or an error in the common link for the reddening and the both metallicity indicators, H$\\beta$. Unfortunately, an error of over 0.03 mag would be required to have a significant impact on the metallicity and, even for this change, would produce [Fe/H] well above the Hyades.  A probable solution to the question is that the indices are, in fact, close to the truth but their interpretation is based upon a questionable extrapolation of the standard relations for stars with more \"normal\" abundances, both in an absolute and a relative sense. This distortion could apply to both the metallicity and the reddening determination, though the size and direction of the effect on the latter variable are unknown. The existence of $uvby$ index distortions due to abundance anomalies has been known for over 20 years with CN variations in CH and Ba stars producing anticorrelated changes in $m_1$ and $c_1$ being a prime example \\citep[see, e.g.,][]{BO80,LB82}. The correlated impact of abundance anomalies on $m_1$, $c_1$, and $hk$ has been documented for field stars \\citep{ATT91,AST} and globular clusters like M22 \\citep{ATC95} and $\\omega$ Cen \\citep{RE00}, but the objects under discussion have been metal-poor giants contained within systems that exhibit star to star variations, either due to inhomogeneities at the time of formation or variations in the evolutionary impact of post-main-sequence internal adjustments. If the source of the anomalous indices in NGC 6253 is an enhanced distribution of $\\alpha$-elements, it is common to all the stars in the cluster and must be the product of galactic chemical evolution. To our knowledge, NGC 6253 is the first open cluster to exhibit potential evidence for a non-scaled-solar element distribution.  How does this solution mesh with observations of the Galactic disk interior to the sun? Two observational pieces of evidence lend some support to this option, though the significance of each remains indeterminate either because the result is too recent to be considered established or because the conclusions remain controversial. One piece of corroborative evidence is supplied by spectroscopic analysis of Cepheids interior to the solar circle by \\citet{AN12}, indicating that Cepheids that lie more than 1.5 kpc closer to the galactic center than the sun have [Fe/H] between +0.2 and +0.4, significantly higher than those found in the solar neighborhood. From simple chemical evolution models, the dramatically higher [Fe/H] should be typical of recently formed stars within this galactocentric radius. It must be emphasized that though the error bars on [Fe/H] are small, so is the sample. If this inner discontinuity is real, it needs to be confirmed with a larger sample extending closer to the galactic center and an explanation for its apparent absence from HII region analyses \\citep{HW99} would be useful.  Another piece of the puzzle is supplied by another well-studied, older open cluster located at a galactocentric distance comparable to NGC 6253 --  NGC 6791, with a controversial history of reddening, metallicity, and age estimation. The most recent comprehensive analysis of the cluster is given by \\citet{CHA}, where CMD comparisons to theoretical isochrones restrict the cluster metallicity to [Fe/H] = +0.4 $\\pm$ 0.1, in excellent agreement with the sole high dispersion spectroscopic abundance analysis of a horizontal branch star by \\citet{PG98}. (For an alternative view of this discussion, see \\citet{TAY}.) If the single spectroscopic result can be trusted as typical of the cluster, it is intriguing to note that the star shows enhancements relative to solar ratios of Mg, Si, Na, and N, but not Ca. Even accepting the high and enhanced metallicity for NGC 6791, its unique properties within the cluster sample, combined with its highly eccentric orbit \\citep{SC95}, have led some to question where it fits within the population mix of the Galaxy, thin disk, thick disk, or bulge. If the last alternative is correct, the cluster is not relevant to understanding the evolution of the disk. If our conjecture regarding the metallicity distribution of NGC 6253 is correct, a trend of high metallicity beginning at least 8 Gyr ago and extending to the present appears established for the disk interior to $\\sim$7 kpc from the galactic center, keeping in mind that the current location for the clusters may not be representative of where they formed."
    },
    "0212/astro-ph0212066_arXiv.txt": {
        "abstract": "We report simultaneous $UBVRI$ photo-polarimetric observations of the long period (7.98 h) AM Her binary V1309 Ori.  The length and shape of the eclipse ingress and egress varies from night to night. We suggest this is due to the variation in the brightness of the accretion stream. By comparing the phases of circular polarization zero-crossovers with previous observations, we confirm that V1309 Ori is well synchronized, and find an upper limit of 0.002 percent for the difference between the spin and orbital periods. We model the polarimetry data using a model consisting of two cyclotron emission regions at almost diametrically opposite locations, and centered at colatitude $\\beta=35^{\\circ}$ and $\\beta=145^{\\circ}$ on the surface of the white dwarf. We also present archive X-ray observations which show that the negatively polarised accretion region is X-ray bright. ",
        "introduction": "The X-ray source RX J0515.6+0105 (V1309 Ori) was preliminarily identified in the {\\sl ROSAT} All Sky Survey as a Cataclysmic Variable (CV) with an orbital period of $\\sim$8 hrs \\cite{b2}. Garnavich et al. \\shortcite{b12} showed that V1309 Ori has properties which are typical of magnetic CVs (mCVs). They also found a deep total eclipse in the light curves. Pointed X-ray observations made using {\\sl ROSAT} \\cite{b45} showed extremely variable X-ray emission in short bursts (up to few seconds), suggesting inhomogeneous accretion of dense blobs.  Detection of variable circular polarization in white light by Buckley \\& Shafter \\shortcite{b4} confirmed the classification of V1309 Ori as an AM Her star (synchronised mCVs).  The orbital period of V1309~Ori 7.98 h is $\\sim$3 hours longer than in any other known AM Her system. Only three AM Hers are known to have periods over 4 hours (RX J1313--32: 4.25 h; AI Tri: 4.59 h; V895 Cen: 4.77 h: Ritter \\& Kolb 1998). The exceptionally long orbital period indicates that the separation between the primary and the secondary is large for mCV.  According to Patterson's \\shortcite{b25} scaling law, magnetic field field strengths up to 150 -- 730 MG should be required to synchronize the spin of the white dwarf with the rotation of the binary system, assuming a mass for the white dwarf $M_{\\rm WD}$=0.6-1.0 M$_{\\odot}$ \\cite{b35}. However, spectroscopic studies have given significantly lower magnetic field values: 33 -- 55 MG Garnavich et al. \\shortcite{b12}; 61 MG Shafter et al. \\shortcite{b35}; \\( < 70 \\) MG de Martino et al. \\shortcite{b8}. Frank, Lasota \\& Chanmugam \\shortcite{b10} have proposed that these low magnetic field values can be understood in terms of the standard evolutionary model, if the system is in a low accretion state for a long enough time or if the magnetic field of the secondary is strong enough (\\( > 1\\)kG) to maintain synchronism.  For these reasons, V1309 Ori is one of the most important objects for studying synchronization and accretion processes in mCVs and their evolution. In this paper we present the results and analysis of our polarimetric, photometric and spectroscopic observations of this system.   \\begin{table*} \\centering \\begin{minipage}{140mm} \\caption{The observing Log for V1309 Ori} \\begin{tabular}{rrrrrrr} \\hline UT Date  & HJD Start & Filter(s) & Duration & Exp. Time & Telescope& Type of Observations \\\\ (at start)  &(2,440,000.0+)& & (hr) & (sec) & & \\\\ \\hline Oct 7 1997 & 10728.5920 & UBVRI & 4.42 & 10 & NOT & photopol. (lin.)\\\\ Oct 8 1997 & 10729.5890 & UBVRI & 4.82 & 10 & NOT & photopol.(cir. \\& lin.)\\\\ Oct 9 1997 & 10730.5875 & UBVRI & 4.65 & 10 & NOT & photopol.(cir. \\& lin.)\\\\ \\hline Nov 24 1997 & 10777.4853 & UBVRI & 8.06 & 10 & NOT & photopol.(cir. \\& lin.)\\\\ \\hline Dec 31 1997 & 10813.6973 & UBVRI & 0.17 & 10 & CASLEO & photopol. (lin.)\\\\ Jan 1 1998 & 10814.5737 & UBVRI & 3.17 & 10 & CASLEO & photopol. (lin.)\\\\ Jan 2 1998 & 10815.5954 & UBVRI & 2.85 & 10 & CASLEO & photopol. (lin.)\\\\ Jan 4 1998 & 10817.5781 & UBVRI & 2.37 & 10 & CASLEO & photopol. (lin.)\\\\ Jan 5 1998 & 10818.5980 & UBVRI & 3.01 & 10 & CASLEO & photopol. (lin.)\\\\ \\hline Dec 25 1998 & 11172.9840 & Blue$^{1}$ \\& Red$^{2}$ & 4.84 & 300 & ANU & spectrosc.\\\\ Dec 26 1998 & 11174.0080 & Blue$^{1}$ \\& Red$^{2}$ & 2.95 & 300 & ANU & spectrosc. \\\\ Dec 27 1998 & 11174.9310 & Blue$^{1}$ \\& Red$^{2}$ & 6.29 & 300 & ANU & spectrosc.\\\\ \\hline \\end{tabular} \\medskip  NOT: 2.56 m Nordic Optical Telescope, Observatorio del Roque de los Muchachos, La Palma, Spain; \\\\ CASLEO: 2.15 m telescope, Complejo Astronomico El Leoncito, province of San Juan, Argentina; \\\\ ANU: 2.30 m Australian National University telescope, Siding Spring Observatory (SSO), Australia; \\\\ Blue$^{1}$: wavelength coverage 3800\\AA -- 5000\\AA\\\\ Red$^{2}$: wavelength coverage 6200\\AA -- 7500\\AA \\\\ \\end{minipage} \\end{table*} ",
        "conclusions": "\\subsection{A well synchronised system}  By comparing our circular polarization curves with those found in the literature we confirm that the spin of the white dwarf in V1309 Ori is synchronised with the orbital period to a high degree. The zero-crossings of circular polarization take place at the same phase of the orbital period as found by Buckley \\& Shafter \\shortcite{b4}, which puts an upper limit of ${\\rm \\sim0.002}$ per cent for the difference of the white dwarf spin and orbital period. There are four polars which have been found to show a small ($\\sim$1 percent) degree of asynchronism. The polar showing the smallest, V1432 Aql, is 0.28 percent asynchronous \\cite{b13}: over two orders of magnitude greater than V1309 Ori.  \\subsection{Accretion geometry}  Our circular polarimetry curves (Figure 2) shows clear positive and negative excursions which indicate that V1309 Ori has two accreting poles. This is consistent with previous polarisation observations. The X-ray data from the {\\sl ROSAT} archive suggests that the negative circularly polarised pole is brighter in X-rays. It is possible that it is bright because of the increased mass transfer at this pole. Alternatively, it might be due to the fact that the accretion flow to this pole is very inhomogeneous, with the dense parts of the flow accreting directly into the white dwarf without causing a shock and therefore liberating its energy at soft X-rays.  We note that Staude et al. \\shortcite{b42} based on optical/UV photometry and optical spectroscopy did not find any evidence for a second accretion pole, although they could not exclude one. Their modelling predicts that the one accreting pole would show a maximum in the soft X-ray light curve at $\\Phi$=0.045 and would show a self eclipse at $\\Phi$=0.55. However, the {\\sl ROSAT} data does not confirm this view. Indeed, we see maximum flux at $\\phi$=0.55 and a minimum at $\\Phi$=0.045.  The X-ray data does show a dip around $\\Phi\\sim$=0.7. This is unlikely to be due to absorption of X-rays by the accretion stream since the stream does not cut through our line of sight to the accretion region at this phase. It is also unlikely that this dip is due to that observation being at a lower accretion state since the same observation shows a peak at $\\phi\\sim$=0.6. The cause of the dip in X-rays is unclear but maybe due to a fraction of the accretion region being self-eclipsed by the white dwarf at these phases.  \\subsection{The mass of the white dwarf} \\label{mass}  Along the ballistic flow, emission extends to ${\\rm V_{x}\\sim-800}$ km s$^{-1}$, ${\\rm V_{y}\\sim-200}$ km s$^{-1}$ in our maps and also the maps of Hoard (1999). The maps of Staude et al.\\shortcite{b42} show the flow extending to ${\\rm V_{y}\\sim-100}$ km s$^{-1}$. Taking $V_{x}, V_{y}$ from our maps (giving a velocity of 820 km s$^{-1}$) and using $v=(2GM_{wd}/r)^{1/2}$ (where $v$ is velocity and $r$ is distance from the white dwarf), we find that the end of the ballistic flow is $r$=2.4 and $2.8\\times10^{10}$ cm distant from the white dwarf for $M_{wd}$=0.6 and 0.7\\Msun respectively.  We can compare these estimates to the expected distance from the white dwarf that material gets coupled by the magnetic field ($R_{\\mu}$) using equation (1b) of Mukai \\shortcite{b23}. We assume $B$=50 MG (typical of the estimates made for V1309 Ori), $\\sigma_{9}$=3 (the radius of the stream in units of $10^{9}$ cm: the value estimated for HU Aqr, Harrop-Allin et al. 1999) and $\\dot{M}_{16}$=10 (the mass transfer rate in units of 10$^{16}$ g s$^{-1}$, Harrop-Allin et al. 1997). We find that for $M_{wd}$=0.6 and 0.7\\Msun we get $R_{\\mu}$= 3.4 and 2.1$\\times10^{10}$ cm respectively. Although there is some considerable degree of uncertainty in how applicable the above formulation for $R_{\\mu}$ actually is, it is interesting that for masses between 0.6--0.7\\Msun the predicted value of $R_{\\mu}$ is consistent with our Doppler maps.  This range of mass is consistent with that estimated by Staude et al. \\shortcite{b42}.  \\subsection{QPOs and their timescales}  Our photometric observations show evidence for QPOs on time scales of 10 min with amplitudes up to 0.2 magnitudes. This compares with 6.7 and 15.5 min: Shafter et al. \\shortcite{b35}. There are some examples of other AM Her systems where flickering on time scales of few minutes have been observed: 4.5 min in BL Hyi \\cite{b39}, 4 -- 11 min in QQ Vul \\cite{b24}, and in AI Tri 6.5-- 7 min and 13.5 -- 14 min \\cite{b36}.  We speculate that the QPOs in V1309 Ori are due to 'blobby accretion', already observed in X-ray data (Walter, Wolk \\& Adams 1995; de Martino et al. 1998).  We compare these timescales with those derived by King \\shortcite{b17}, and King \\shortcite{b18} (see also Chanmugam 1995). The irradiation of the accretion flow may ionise the subsonic accretion flow below the inner $L_{1}$ point and modulate gas flow through this point on the timescale of the dynamical time scale in the Roche potential near $L_{1}$.  The equation presented by King \\shortcite{b17}  \\begin{equation} ~~~~~~~~~~~~~~~~~T_{osc}\\sim\\frac{H_{*}}{c_{*}}=5.5\\times10^{-2}P \\end{equation}  where ${\\rm H_{*}}$ is the scale height and ${\\rm c_{*}}$ is local sound speed near ${\\rm L_{1}}$, predicts the timescales for these oscillations. The orbital period of 479 minutes gives us a prediction of 26 minutes for V1309 Ori. The observed timescales of QPOs in V1309 Ori, 10 and 15 minutes, are approximately half of the predicted. The QPOs are seen strongest in $UBV$, and are negligible in the longer wavelengths, which may due to reason that only a fraction of flow will undergo these oscillations near ${\\rm L_{1}}$, as pointed out by King \\shortcite{b18}.  \\subsection{Eclipse profiles}  Several eclipsing polars have been observed with high signal to noise and high time resolution. These include HU Aqr \\cite{b15} and UZ For \\cite{b26}. Both of these systems show a sharp eclipse ingress lasting several seconds. This sharp drop in intensity is associated with the eclipse of the bright accretion region on the white dwarf. In the case of UZ For there are two sharp intensity drops, indicating that there are two accretion regions visible during the eclipse.  One of the most striking features about the eclipse profiles of V1309 Ori is the obvious lack of a sharp ingress or egress which indicate the (dis)appearance of the white dwarf and/or hot spots in the surface of the white dwarf behind the secondary. The relative faintness of the accretion region(s), compared to the bright accretion stream in V1309 Ori, may be the reason for this.  After the eclipse of the white dwarf, the accretion stream is still visible for a length of time. Observations of other polars such as HU Aqr (Glenn et al. 1994, Bridge et al. 2002) show that this length of time can vary from one cycle to the next. The fact that we observe a variable eclipse ingress of the accretion stream in V1309 Ori is therefore not un-typical of polars. However, what does make V1309 Ori unique amongst polars is the fact that the eclipse egress is highly variable: all other polars show a rapid rise at the same phase coming out from eclipse. The fact that V1309 Ori does not implies either that we can observe the stream above the orbital plane before the white dwarf is visible, or that the accretion stream travels far enough around the white dwarf so that it is visible before the white dwarf itself.  To investigate if this further, we shown in Figure \\ref{system} the view of the system at two different phases ($\\Phi$= 0.96 and 1.04). We use the following system parameters in determining these: $i=78^{\\circ}$, $q$=0.67, $M_{wd}=0.7$ \\Msun~\\cite{b42} and a white dwarf - secondary star separation of 1.47$\\times10^{11}$ cm (determined using the above parameters and standard Roche lobe geometry). We also show a single magnetic field line originating from the negative circularly polarised accretion region ($\\beta=35^{\\circ}$ and face onto the observer at $\\Phi$=0.2, \\S \\ref{modelling}). In Figure \\ref{system}, the accretion streams leading to both poles are visible at $\\rm \\Phi=0.96$. At $\\rm \\Phi=0.04$ only the stream leading to the negative pole is visible. The white dwarf appears before the stream leading to the positive pole is visible. Even if the negative pole is not visible at 0.04, the emission from the stream leading to that pole is. We take $R_{\\mu}/a$=0.2: this implies $R_{\\mu}=2.9\\times10^{10}$ cm. This is consistent with our findings in \\S \\ref{mass}. This shows that for this accretion stream geometry the accretion stream is visible after the white dwarf has been eclipsed and also before the white dwarf comes out of eclipse. If the stream emission was highly variable then this could explain the variable egress profile.  \\begin{figure*} \\begin{center} \\setlength{\\unitlength}{1cm} \\begin{picture}(8,6) \\put(-3,-2.){\\special{psfile=mb905fig11.eps hscale=34 vscale=34}} \\put(3,-2.){\\special{psfile=mb905fig12.eps hscale=34 vscale=34}} \\end{picture} \\end{center} \\caption{The view of the system at phases $\\Phi$=0.96 and 1.04. See the text for details of the system parameters and accretion geometry which are assumed. The secondary Roche lobe is projected as a wire model so that the white dwarf and stream remain visible.} \\label{system} \\end{figure*}  \\subsection{Evolution}  Garnavich et al. \\shortcite{b12} and Shafter et al. \\shortcite{b35} noted that the secondary star in V1309 Ori is oversized for its spectral type (M0 - M1) and mass (0.4--0.6 $M_{\\odot}$). Indeed, recent binary evolution models (eg Smith \\& Dhillon 1998 and Baraffe \\& Kolb 2000) which assume an unevolved donor star and typical mass transfer rates, all predict either a much earlier or later spectral type for V1309 Ori than observed (cf Figure 1 of Baraffe \\& Kolb 2000).  However, by assuming an evolved donor it is possible to match the observed spectral type to the predicted value. For instance, from Figure 3 of Baraffe \\& Kolb \\shortcite{b1}, for an initial secondary star mass of $M_{2}$=1.2\\Msun, a central Hydrogen abundance at the start of mass transfer of $X_{c}$=0.05, and a mass transfer rate of $\\dot{M} =1.5 \\times 10^{-9} M_{\\odot} (\\sim1\\times10^{17}$) g s$^{-1}$) we find we that a spectral type of M1 is predicted for an orbital period of 8 hrs. This is within the estimated range of values required to satisfy conditions for observed oversized secondary to fill its Roche lobe (the main uncertainties are the distance and the mass of the white dwarf, eg  Harrop-Allin et al. 1997, de Martino et al. 1998).  As already noted by King, Osborne \\& Schenker \\shortcite{b19}, V1309 Ori is indeed a good possible candidate for a binary system which has gone through a supersoft source phase and contains a nuclear evolved donor star. Szkody \\& Silber \\shortcite{b43} and Schmidt \\& Stockman \\shortcite{b33} have noticed that V1309 Ori has extraordinary strong excitation lines of ${ \\rm N~V~\\lambda 1240}$ and ${ \\rm O~V~\\lambda1370}$) which may support such an interpretation."
    },
    "0212/astro-ph0212316_arXiv.txt": {
        "abstract": "The rapidly growing base of observational data for supernova explosions of massive stars demands theoretical explanations. Central of these is a self-consistent model for the physical mechanism that provides the energy to start and drive the disruption of the star. We give arguments why the delayed neutrino-heating mechanism should still be regarded as the standard paradigm to explain most explosions of massive stars and show how large-scale and even global asymmetries can result as a natural consequence of convective overturn in the neutrino-heating region behind the supernova shock. Since the explosion is a threshold phenomenon and depends sensitively on the efficiency of the energy transfer by neutrinos, even relatively minor differences in numerical simulations can matter on the secular timescale of the delayed mechanism. To enhance this point, we present some results of recent one- and two-dimensional computations, which we have performed with a Boltzmann solver for the neutrino transport and a state-of-the-art description of neutrino-matter interactions. Although our most complete models fail to explode, the simulations demonstrate that one is encouragingly close to the critical threshold because a modest variation of the neutrino transport in combination with postshock convection leads to a weak neutrino-driven explosion with properties that fulfill important requirements from observations. ",
        "introduction": "\\label{sec:problem}  The primary energy source for powering supernovae of massive stars is the gravitational binding energy of the newly formed proto-neutron star or proto-black hole (energy from nuclear reactions contributes at a minor level). To initiate and drive the explosion, energy from some temporary storage, e.g., internal or rotational energy of the compact remnant, must be transferred to the outer stellar layers to be finally converted to kinetic energy of the ejecta. This might be achieved by hydrodynamical shocks, by neutrinos, or by magnetic fields as mediators. Accordingly, one distinguishes between \\begin{itemize} \\item[(i)] the (prompt) mechanism, which works on a dynamical timescale by the hydrodynamical shock that is created at the moment of core bounce, \\item[(ii)] the delayed, neutrino-driven mechanism which starts the explosion on the secular timescale of neutrino-energy deposition behind the supernova shock, \\item[(iii)] and the magnetohydrodynamical (MHD) mechanism, which requires that initial seed magnetic fields are amplified to a dynamically relevant strength by differential rotation. \\end{itemize}  Intense radiation or relativistic outflows of charged particles might also play a role for very special conditions. They may, for example, originate from the vicinity of an accreting black hole that has formed after the collapse of a rotating stellar core. Relativistic jets as driving mechanism are currently discussed for stellar explosions that have been observed in association with gamma-ray bursts (see the contributions by S.~Woosley and A.~MacFadyen at this conference).  Depending on the mediator, the conditions for efficient energy transfer, the corresponding timescale, and the tapped energy reservoir are different. A lot of work has been spent in the past 40 years on the search for a viable supernova mechanism and the study of the various theoretical suggestions. A brief review of these efforts and the current status of our knowledge can be found in Ref.~\\cite{janbur02b}.  The relevance of the different mechanisms for stellar explosions listed above depends on the (poorly known) physical conditions in collapsed stellar cores and on the properties of the progenitor stars. Some of the involved requirements are more likely to be fulfilled than others, some combinations of necessary conditions may be more frequent and more typical, while others may be realized only in rare cases and for very special, exceptional circumstances.  The neutrino-driven mechanism~\\cite{wil85,betwil85} involves a minimum of controversial assumptions and uncertain degrees of freedom in the physics of collapsing stars. It relies on the importance of neutrinos and their energetic dominance in the supernova core. After the detection of neutrinos in connection with Supernova~1987A and the overall confirmation of theoretical expectations for the neutrino emission, this can no longer be considered as a speculative assumption but is an established fact. Of course, this does not mean that such a minimal input is sufficient to understand the cause of supernova explosions and to explain all observable properties of supernovae. But at least it can be taken as a good reason to investigate how far one can advance with a minimum of imponderabilities. ",
        "conclusions": "\\label{sec:conclusions}   Our 2D models with a Boltzmann solver for the neutrino transport have considerably reduced the uncertainties associated with the treatment of the neutrino physics in previous multi-dimensional simulations. With the most complete implementation of the transport physics we could not obtain explosions. This result suggests that the neutrino-driven mechanism fails with the employed input physics, at least in case of the considered 15$\\,$M$_{\\odot}$ star. We do not think that the remaining uncertainties in our simulations (mainly the approximate treatment of general relativistic effects) are likely to jeopardize this conclusion. A comparison with fully relativistic one-dimensional calculations (Liebend\\\"orfer, personal communication~\\cite{lieetal02}) is very encouraging. Because of the remarkable similarity of the shock trajectories of different progenitors in spherical symmetry~\\cite{liemes02}, it is likely that our negative conclusion is also valid for other pre-collapse configurations with a similar structure. Significant star-to-star variations of the progenitor properties with a non-monotonic dependence on the stellar mass~\\cite{wooheg02}, however, suggest that multi-dimensional core-collapse simulations of a larger sample of progenitors are needed before one can make final, more generally valid statements. The supernova problem is highly nonlinear and surprises may lurk behind every corner.  It would therefore be premature to conclude that the neutrino-driven mechanism fails and that not even postshock convection can alter this unquestioned outcome of all current spherical models. Besides studying other progenitors with multi-dimensional simulations, one should also investigate the effects of rotation and the influence of different high-density equations of state on the long-time post-bounce evolution and the neutrino-heating phase in a supernova. There is considerable uncertainty associated with the poorly known physics in the nuclear and supranuclear medium.  Our successfully exploding 2D model, Model~s15Gio\\_2d.a, at least demonstrates that simulations which include the effects of postshock convection are rather close to an explosion. Therefore modest changes of the neutrino emission and transport seem to be already sufficient to push them beyond the critical threshold. The properties of the explosion in this case are very encouraging and may support one's belief in the basic viability of the delayed explosion mechanism. At 380$\\,$ms after bounce the shock has arrived at a radius of more than 2500$\\,$km and is expanding with about 10000$\\,$km/s. The explosion of this model does not seem to become very energetic. It is only $\\sim 4\\times 10^{50}\\,$erg at that time, but still increasing. This may not be a serious problem if one recalls the large spread of energies of observed supernovae (Supernova~1999br, for example, is estimated to have an ejecta mass of $14\\,$M$_{\\odot}$ and an explosion energy of about $6\\times 10^{50}\\,$erg~\\cite{ham02}).  Since the explosion starts rather late (at $\\sim 150\\,$ms post bounce), the proto-neutron star has accreted enough matter to have attained an initial baryonic mass of 1.4$\\,$M$_{\\odot}$. Therefore our simulation does not exhibit the problem of previous successful multi-dimensional calculations which produced neutron stars with masses on the lower side of plausible values ($\\sim 1.1\\,$M$_{\\odot}$). Also another problem of published explosion models (e.g.,~\\cite{herben94,burhay95,janmue96,fry99}) has disappeared: The ejecta mass with $Y_e \\la 0.47$ is less than $10^{-4}\\,$M$_{\\odot}$, thus fulfilling a constraint pointed out by Hoffman et.~al.~\\cite{hofwoo96} for supernovae if they should not overproduce the $N=50$ (closed neutron shell) nuclei, in particular $^{88}$Sr, $^{89}$Y and $^{90}$Zr, relative to the Galactic abundances. Of course, final statements about explosion energy, ejecta composition, and the neutron star mass (which may grow by later fallback, especially when the explosion energy remains low) require to follow the explosion for a longer time.    \\bigskip\\noindent {\\small {\\bf Acknowledgements:} We are grateful to K.~Takahashi for providing routines to calculate the improved neutrino-nucleon interactions, and to C.~Horowitz for correction formulae for the weak magnetism. We also thank M.~Liebend\\\"orfer for making output data of his simulations available to us for comparisons. The Institute for Nuclear Theory at the University of Washington is acknowledged for its hospitality and the Department of Energy for support during a visit of the Summer Program on Neutron Stars, during which most of the work leading to Fig.~\\ref{fig:monomode} was done. HTJ, RB and MR are grateful for support by the Sonderforschungsbereich 375 on ``Astroparticle Physics'' of the Deutsche Forschungsgemeinschaft. TP was supported in part by the US Department of Energy under Grant No. B341495 to the Center of Astrophysical Thermonuclear Flashes at the University of Chicago, and in part by the grant 2.P03D.014.19 from the Polish Committee for Scientific Research. He performed his simulations on the CRAY SV1-1A at the Interdisciplinary Centre for Computational Modelling in Warsaw. The 2D simulations with Boltzmann neutrino transport were only possible because a node of the new IBM ``Regatta'' supercomputer was dedicated to this project by the Rechenzentrum Garching. Computations were also done on the NEC SX-5/3C of the Rechenzentrum Garching, and on the CRAY T90 and CRAY SV1ex of the John von Neumann Institute for Computing (NIC) in J\\\"ulich. }"
    },
    "0212/astro-ph0212120_arXiv.txt": {
        "abstract": "We examined the kinematics, ionization conditions, and physical size of the absorption clouds in a $z=1.3911$ damped {\\Lya} absorber (DLA) in the double image lensed quasar Q~$0957+561$ A,B (separation $135h_{75}^{-1}$~pc at the absorber redshift).  Using HIRES/Keck spectra (FWHM$~\\simeq 6.6$~{\\kms}), we studied the {\\MgIIdblt} doublet, {\\FeII} multiplet, and {\\MgI} $\\lambda 2853$ transition in absorption.  Based upon the {\\FeII} profiles (the {\\MgII} suffers from saturation), we defined six ``clouds'' in the system of sightline A and seven clouds in system of sightline B.  An examination of the $N(v)$ profiles, using the apparent optical depth method, reveals no clear physical connection between the clouds in A and those in B.  The observed column density ratios of all clouds is $\\log N({\\MgI})/N({\\FeII}) \\simeq -2$ across the full $\\sim 300$~{\\kms} velocity range in both systems and also spatially (in both sightlines).  This is a remarkable uniformity not seen in Lyman limit systems.  The uniformity of the cloud properties suggests that the multiple clouds are not part of a ``halo''.  Based upon photoionization modeling, using the $N({\\MgI})/N({\\FeII})$ ratio in each cloud, we constrain the ionization parameters in the range $-6.2 \\leq \\log U \\leq -5.1$, where the range brackets known abundance ratio and dust depletion patterns.  The inferred cloud properties are densities of $2 \\leq n_{\\rm H} \\leq 20$~{\\cc}, and line of sight sizes of $1 \\leq D \\leq 25$~pc.  The masses of the clouds in system A are $10 \\leq M/M_{\\odot} \\leq 1000$ and in system B are $1 \\leq M/M_{\\odot} \\leq 60$ for spherical clouds.  For planar clouds, the upper limits are $400h_{75}^{-2}$~M$_{\\odot}$ and $160h_{75}^{-2}$~M$_{\\odot}$ for A and B, respectively.  We favor a model of the absorber in which the DLA region itself is a single cloud in this complex, which could be a parcel of gas in a galactic ISM.  We cannot discern if the {\\HI} in this DLA cloud is in a single, cold phase or in cold+warm phases.  A spherical cloud of $\\sim 10$~pc would be limited to one of the sightlines (A) and imply a covering factor less than 0.1 for the DLA complex.  We infer that the DLA cloud properties are consistent with those of lower density, cold clouds in the Galactic interstellar medium. ",
        "introduction": "\\label{sec:intro}  Damped {\\Lya} absorbers (DLAs) are well appreciated as useful astronomical laboratories for measuring metal enrichment, nucleosynthetic processing, and dust evolution from $z=4$ to the present epoch \\citep[e.g.][]{pettini94,lu96,pettini97,vladilo98,pettini99,pw00,prochaska02,ledoux02}. Their redshift number density evolution \\citep{wolfe95,lisa96,rt00} places constraints on $\\Omega ({\\HI},z)$, and probes the {\\HI} mass function for column densities greater than $N({\\HI}) = 2\\times 10^{20}$~{\\cmsq}.  The gas kinematics measured from low--ionization metal absorption lines in DLAs provides valuable data upon which hypotheses for their dynamical formation, evolution, and physical nature can be based.  To date, these competing hypotheses include thick rotating disks \\citep{pw97,pw98}, merging cold dark matter halos \\citep{haehnelt98,mcdonald99}, and supernovae winds \\citep[][however, see Bond et~al.\\ 2001]{nulsen98,shaye01}.  The fact that DLAs provide leverage on questions relating to cosmological chemical evolution, nucleosynthetic and dust evolution, galaxy and/or hierarchical clustering evolution, and global star formation evolution makes them one of the most useful, and therefore, important astronomical objects of study.  Of the great deal that has been learned about DLAs, one property remains elusive; there is little direct observational evidence from which the characteristic size of the DLA ``region'' itself can be deduced.  By DLA region, we mean the parcel of gas that is associated with the very large {\\HI} column density greater than $N({\\HI}) = 2 \\times 10^{20}$~{\\cmsq}.  From the size estimates, one could then deduce the covering factor of DLAs within their host objects, as well as their densities, and most importantly, their masses (though this quantity is geometry dependent).  Associating DLAs with the normal, bright galaxy population and assuming a unity covering factor, \\citet{steidel93} deduced a characteristic size\\footnote{Throughout this paper we use $h_{75} = H_{0}/75$ and $q_{0}=0$.} of $\\sim 20h^{-1}_{75}$~kpc.  However, it is now known that DLAs are associated with galaxies having a wide variety of morphological types, including $0.1~L^{\\ast}$ galaxies and low surface brightness (LSB) galaxies \\citep{lebrun97,rt98,bowen01}. Indeed, the luminous host of many DLAs remain undetected even after dedicated searches to find their diffuse {\\Lya} emission \\citep[e.g.,][]{deharveng90,lowenthal95}, their optical continuum emission \\citep[e.g.,][]{steidel97,kulkarni00,kulkarni01} and their {\\Ha} emission \\citep[e.g.,][]{bunker99,bouche00,kulkarni00,kulkarni01}.  In the absence of a well--defined luminosity function for the luminous hosts of DLAs, and without knowledge of the covering factor of the DLA region in these objects, we remain ignorant of a characteristic DLA size.  It is safe to say, however, that estimates incorporating low surface brightness galaxies in the luminosity function of known DLA hosts would drive down the characteristic size.  Scaling the Steidel estimate by the ratio of $\\Phi _{\\ast}$ for the luminosity functions of normal bright galaxies \\citep{lilly95,ellis96,lin99} and LSB galaxies \\cite[e.g.,][]{dalcanton97}, we find an approximate characteristic size of $(0.033/(0.033+0.08)\\cdot 20h^{-1}_{75} \\sim 6h_{75}^{-1}$~kpc.  The data directly constraining the physical sizes of DLAs are few.  In the lensed quasar HE~$1104-1805$, \\citet{smette95} placed limits of $5$--$17h^{-1}_{75}$~kpc on a DLA in one of the images.  A more constraining result was obtained by \\citet{michal97}, who used the lensed quasar Q~$0957+561$ A,B to show that the physical extent of a DLA at $z=1.3911$ is less than $135h_{75}^{-1}$~pc.  On the other hand, \\citet{petitjean00} suggest a lower limit of $200h_{75}^{-1}$~pc for a probable DLA in the lensed quasar APM $08279+5255$ (based upon partial covering arguments).  Using VLBA observations of the slightly extended radio quasar B~$0738+313$, \\citet{lane00} find $\\log N({\\HI}) = 21.2$~{\\cmsq} in 21--cm absorption at two locations separated by $20h^{-1}_{75}$ pc [we note, however, that 21--cm emission line techniques may not be selecting objects that are one and the same as the DLAs selected using quasar absorption line techniques \\citep[see][]{rt00,cwc01}].  In addition to the great insights garnished from low-- to moderate--resolution spectroscopic studies of multiply imaged quasars, high resolution spectra provide constraints on the kinematics of metal--line systems at the few kilometers per second level \\citep[e.g.,][]{rauch99-i,lopez99,petitjean00,rauch02-iv}.  This means the gaseous structure can be studied component by component, directly providing their sightline to sightline velocity shears, $\\Delta v/\\Delta r$, where $r$ is the physical separation of the sightlines at the absorber.  Such measurements allow insights into the ionization structures, masses, density gradients, and energy input budgets \\citep[e.g.,][]{rauch99-i,rauch01-iii,dodorico02}.  It is important to realize that these insights are derived based upon models of the gas dynamics and geometry.  These models, however, are strongly founded upon observations of the gaseous components in the Milky Way, LMC, SMC, and other local galaxies.  The reasonably bright gravitationally lensed quasar Q~$0957+561$ A,B at $z_{em} = 1.41$ with separation 6{\\arcsec} provides an excellent opportunity to study high--resolution spectra of two sightlines through a $z>1$ DLA \\citep[e.g.,][]{walsh79,wills80,young80,young81a,young81b,turnshek93,michal97,zuo97,pettini99}. The $z=1.3911$ absorber has a neutral hydrogen column density of $N({\\HI}) = 1.9\\pm 0.3\\times 10^{20}$~{\\cmsq} in the A spectrum and $N({\\HI})=8\\pm2 \\times 10^{19}$~{\\cmsq} in the B spectrum \\citep{zuo97}.  This suggests that the DLA region is concentrated in front of the A image.  At the redshift of the absorber, the physical separation of the A and B sightlines is $135h_{75}^{-1}$~pc \\citep{smette92}.  \\citet{zuo97} found a differential reddening in spectrum A as compared to spectrum B and deduce a dust to gas ratio of $\\sim 0.6$ in the DLA.  Imaging of the quasar field has been reported by \\citet[][ground based]{dalhe94} and \\citet[][WFPC2/HST]{bernstein97}.  Despite these efforts, the luminous object associated with the $z=1.3911$ absorber remains unidentified.  The most extensive study of the DLA to date, using FOS/{\\it HST} spectra ($R=1300$), was presented by \\citet{michal97}.  In general, they found ``the differences in the damped Lyman series absorption in the lensed components are the only significant spectral characteristics that distinguishes the far--ultraviolet spectra of Q~$0957+561$ A,B.''  Assuming a spherical geometry and a neutral hydrogen fraction of $N({\\HI})/N({\\rm H}) = 0.05$ \\citep{chartas95}, the authors estimate the DLA mass to be $2.6\\times 10^{6}$~M$_ {\\odot}$, which they interpreted to be consistent with a giant molecular cloud or a small condensation (spiral arm) within a galactic disk (viewed pole on).  Because the metal lines are tightly correlated for the two sightlines, \\citet{michal97} suggested that the large metal--line absorption arises in a galactic ``halo''.  In this paper, we present $R=45,000$ HIRES/Keck spectra of the Q~$0957+561$ A,B.  We focus on the {\\MgIIdblt} doublet, the {\\MgI}~$\\lambda 2853$ transition, and several {\\FeII} transitions arising in the DLA at $z=1.3911$.  In \\S~\\ref{sec:data}, we present the data, and briefly describe the data reduction and analysis.  We present the observed kinematic and spatial properties of the absorber column densities in \\S~\\ref{sec:cols}.  Assuming photoionization equilibrium, we model the clouds and present their densities, sizes, and masses in \\S~\\ref{sec:results}.  We briefly discuss these properties in \\S~\\ref{sec:discuss} and provide concluding remarks in \\S~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have observed the images of the lensed quasar Q~$0957+561$ A,B with the HIRES/Keck--I instrument (resolution FWHM$~\\simeq 6.6$~{\\kms}). We have presented an analysis of the {\\MgII}, {\\MgI}, and {\\FeII} absorption profiles from the $z=1.3911$ DLA system.  We adopted the hydrogen column densities for systems A and B from \\citet{zuo97}, which are $\\log N({\\HI}) = 20.3$ and 19.9~{\\cmsq}, respectively.  The line of sight separation is $\\simeq 135h^{-1}_{75}$~pc at the redshift of the absorber \\citep{smette92}.  We converted the absorption profiles to their apparent optical depth column density \\citep[$N(v)$,][]{savage91}.  Based upon the location of local minima in the $N(v)$ profiles for {\\FeII}, we defined six ``clouds'' in system A and seven clouds in system B and integrated the $N(v)$ to obtain the cloud column densities.  There is a ``dominant'' cloud in each line of sight.  It may be that these clouds contain the bulk of the neutral hydrogen gas.  If the cloud geometry is planar and extents across sightline B, then there is a neutral hydrogen column density gradient of $9\\times 10^{17}h_{75}$~{\\cmsq}~pc$^{-1}$ and a velocity sheer of $\\simeq 0.2h_{75}$~{\\kms}~pc$^{-1}$.  The clouds were assumed to be in photoionization equilibrium.  Using Cloudy \\citep{ferland}, we modeled the clouds as constant density, plane--parallel ``slabs'' illuminated on one face by the ultraviolet background ionizing spectrum.  We used the $N({\\MgI})/N({\\FeII})$ ratio in each cloud to constrain the ionization conditions.  Since both Mg and Fe suffer dust depletion and originate predominantly in separate nucleosynthetic environments, we bracketed the [Mg/Fe] abundance pattern for the range of dust depletions seen in the Galaxy and LMC/SMC and for the observed abundance patterns in the local universe.  The observed $N({\\MgI})/N({\\FeII})$ ratio is remarkably uniform.  Not only are the ratios consistent with $\\log [N({\\MgI})/N({\\FeII})] = -2$ across the full velocity range in both systems, but they are also consistent with this value spatially (in both sightlines).  This resulted in very uniform cloud physical properties as inferred from the photoionization modeling.  The ionization parameter of the clouds is in the range $-6.2 \\leq \\log U \\leq -5.1$.  This yielded clouds with densities of $2 \\leq n_{\\rm H} \\leq 20$~{\\cc} and line of sight physical extents of $1 \\leq D \\leq 25$~pc.  The inferred masses are geometry dependent.  For spherical geometries the masses of the clouds in system A are $10 \\leq M/M_{\\odot} \\leq 1000$ and in system B are $1 \\leq M/M_{\\odot} \\leq 60$.  For cylindrical geometries constrained by the line--of--sight separation of less than 200~pc, the cloud masses have upper limits of $400h_{75}^{-2}$~M$_{\\odot}$ and $160h_{75}^{-2}$~M$_{\\odot}$ for systems A and B, respectively.  These cloud properties are consistent with those for lower density, cold clouds in the Galactic interstellar medium \\citep{spitzer85,savage96}.  We focused our discussion on the physical nature of the DLA ``region'', the object that actually gives rise to the damped {\\Lya} absorption of $\\log N({\\HI}) = 20.3$~{\\cmsq}.  Based upon simulations, we favor a picture in which the DLA is a single cloud in the multi--cloud profiles.  We cannot discern, however, if the DLA comprises a ``cold'' single ionization phase, as suggested by our photoionization models, or a ``cold+warm'' two--phase gas complex.  If the DLA cloud is spherical in nature, then its size is on the order of $\\sim 10$~pc, and it is limited to one of the sightlines (A).  This implies a covering factor of less than 0.1.  The other multiple gas clouds in the proximity of this small DLA cloud would have to have experienced the same sources of nucleosynthetic enrichment, be optically thick in $N({\\HI})$, and have similar dust contents.  This implies that the material distributed in proximity to the DLA is well mixed and ionized uniformly.  This is in stark contrast to the significant variations seen in Lyman limit systems \\citep[e.g.,][]{cv01}, which are thought to arise in the outer disks and halos of galaxies.  As such, we suggest that the low ionization clouds accompanying DLAs are not arising in galactic halos.  Rather, we infer that DLAs arise in small gas--rich regions within galaxies.  The data and models suggest that these regions are complexes comprised of small, optically thick clouds similar to the lowest mass, cold {\\HI} clouds in the Galaxy.  Furthermore, the data suggest that they are well mixed chemically and have similar photoionization conditions."
    },
    "0212/astro-ph0212470_arXiv.txt": {
        "abstract": "We will present a multiepoch study of the three blazars 0954+658 (BL--Lac), PKS\\,1510--089 (HPQ) and 1749+096 (BL--Lac). The first two sources are known to be $\\gamma$--ray loud. Our study is based on milliarcsecond resolution polarimetric observations carried out with the VLBA at 8.4 GHz. The observations took place between January 1999 and May 2001. Superluminal motion is detected along the jet of PKS\\,1510--089 and 1749+096, with $\\beta_{app} \\sim 10$ for all features. Magnetic field structure is revealed along the jets of 0954+658 and PKS\\,1510--089. The polarisation properties of the parsec--scale jets remain stationary in all sources, regardless of their total flux density variability in the radio band and of the presence of superluminal features. ",
        "introduction": "\\label{sec1}  The blazar phenomenon is now believed to be the result of an anisotropic relativistic jet propagating at a small angle to the observer's line of sight. Variability at all wavelengths and the peculiar radio properties on the parsec scale, such as one--sidedness, jet bending, misalignement and superluminal features, can all be accounted for under these assumptions.\\\\ In a number of blazars the emission extends out to the MeV regime, i.e. to the $\\gamma$--rays. Many models have been proposed to explain the origin of this very high energy emission, such as for example synchrotron self-Compton (i.e. Maraschi et al. 1992), inverse Compton scattering, with a variety of origins for the scattered photons (i.e. Dermer et al. 1992; Sikora et al. 1994), synchrotron emission by ultra relativistic electrons and positrons (Ghisellini et al. 1993). However, no consensus on the dominant process has been reached yet.\\\\ The connection between the $\\gamma$--ray emission and the parsec--scale radio properties is under study, and many questions are still unexplained. For example, it is still unclear if (1) $\\gamma$--ray loud blazars are characterised by higher apparent speeds and/or higher Lorentz factors; (2) jet bending, wiggling and misalignement are more prominent in $\\gamma$--ray loud sources. Beyond this, the correlation between radio variability and changes in the magnetic field structure is still unclear for all classes of radio loud AGNs.\\\\ In order to increase the statistics, and to shed light on these important questions, detailed multifrequency studies for a high number of radio loud blazars is necessary. In this paper we present a multiepoch study of three radio blazars, carried out at 8.4 GHz with the Very Long Baseline Array (VLBA). Among our sources, two are also $\\gamma$--ray loud.\\\\ We defined $h$=H$_{0}$/100 and $q_{0}=0.5$. ",
        "conclusions": "\\label{sec4} Our monitoring program on three blazars revealed different properties on the parsec--scale jets. In particular:  \\noindent -- superluminal motion is detected in PKS\\,1510--089 (two superluminal knots) and in 1749+096. In all cases an apparent velocity $\\beta_{app} \\sim 10~h^{-1}$c is found;  \\noindent -- a stationary jet is present in 0954+658, despite the major variability in the radio band during the period of our monitoring program;  \\noindent -- the polarisation properties along the parsec--scale jet remain stationary from epoch to epoch in all sources, regardless of their flux density variability and the presence of superluminal features;  \\noindent -- magnetic field structure is found in PKS\\,1510--089, where a higher percentage of polarisation is detected along the outer edges of the jet."
    },
    "0212/astro-ph0212193_arXiv.txt": {
        "abstract": "{I present in this paper a method to subtract the bias due to source photon noise from visibilities measured with a single-mode optical interferometer. Properties of the processed noise are demonstrated and examples of subtraction on real data are presented. ",
        "introduction": "The properties of source photon noise are well known.  It follows a Poisson distribution whose variance is equal to the average total number of photons.  In frequency space, it is a white noise with a flat average power spectral density.  In most practical cases where the observables are linearly linked to the number of photons detected, photon noise can be directly averaged out from the data to reduce its variance.  For some applications for which the observables are quadratically linked to the number of photons, the data suffer both from photon noise and from a bias linked to the variance of the noise. This is for example the case in speckle imaging techniques where the source spatial intensity distribution is recovered from the power spectral density of a short time exposure (\\cite{thiebault1994}).  In astronomical optical interferometry, the observables are the modulus of the visibility and its phase usually expressed as a closure phase quantity.  The visibility modulus can be obtained by integrating the modulus of the spectrum of interferograms.  However, this estimator is biased by the power spectral density of noises as these add to the power spectral density of the fringe signal.  An unbiased estimator of the modulus of the visibility is obtained by forming the squared modulus of the visibility as the power spectral densities of the noises can be independently estimated and subtracted. An example in interferometry is the computation of the fringe squared visibility in the ABCD method where the white light fringe is sampled at four $\\lambda/4$ spaced optical path differences. An unbiased single fringe ABCD estimator is obtained by subtracting the source photon noise and the detector noise variances (when the noise has a flat spectrum the power spectral density is constant and equal to the variance) from the fringe power spectral density (\\cite{tango1980}).  \\\\ In single-mode interferometers, beams are spatially filtered by single-mode waveguides trading phase fluctuations against intensity fluctuations.  A fraction of intensities collected by each aperture can be measured to renormalize interferograms to eliminate the fluctuations due to turbulence.  The visibility estimator is no longer directly linked to the power spectral density of the fringe signal.  I demonstrate in the following sections that the classical method (explained further in the paper) established by \\cite{goodman1985} can be rigourously extended to such ratios of physical noisy signals under certain assumptions to provide unbiased visibility estimators.  Real data reduction cases are presented to illustrate the method. ",
        "conclusions": "A method to subtract the photon noise bias from visibility data acquired with a single-mode interferometer has been presented. An analytical expression for the bias can be established under verified assumptions. Other non-analytical methods based on the fit of the average level of power spectral densities may suffer from confusion with fringe signal and from detector instabilities and the analytical method should be preferred.  \\begin{acknowledgement} I wish to thank P. Bord\\'e for his careful reading of the paper and for his precious comments that helped me write the final version. \\end{acknowledgement}"
    },
    "0212/hep-ph0212379_arXiv.txt": {
        "abstract": "\\vspace{1cm} The relic abundance and the scalar cross--section off nucleon for light neutralinos (of mass below about 45 GeV) are evaluated in an effective MSSM model without GUT--inspired relations among gaugino masses. It is shown that these neutralinos may provide a sizeable contribution to the matter density in the Universe and produce measurable effects in WIMP direct detection experiments. These properties are elucidated in terms of simple analytical arguments. ",
        "introduction": "\\label{sec:intro}  Most works on relic neutralinos consider supersymmetric schemes with a unification assumption for the gaugino masses $M_i$ ($i = 1,2,3$) at the GUT scale $M_{GUT} \\sim 10^{16}$ GeV. This hypothesis implies that at lower scales the following relations hold: \\begin{equation} M_1 : M_2 : M_3 = \\alpha_1 : \\alpha_2 : \\alpha_3, \\label{sec:unif} \\end{equation} where the $\\alpha_i$ ($i = 1,2,3$) are the coupling constants of the three Standard Model gauge groups. In particular, at the electroweak scale, $M_{EW} \\sim 100$ GeV, $M_1$ and $M_2$ are related by the expression: \\begin{equation} M_1 =  \\frac{5}{3} \\tan^2 \\theta_{W} \\; M_2 \\simeq 0.5~M_2. \\label{eq:gut} \\end{equation} However, there are theoretical arguments for considering supersymmetric schemes where the unification assumption on gaugino masses is not satisfied\\cite{nonunif}.  In the present paper we analyse properties of relic neutralinos in an effective Minimal Supersymmetric extension of the Standard Model (MSSM) where the GUT relation of Eq. (\\ref{eq:gut}) is relaxed. Previous papers where supersymmetric schemes without gaugino masses unification have been considered in connection with relic neutralinos include the ones reported in Refs. \\cite{dt,gr,mny,go,bbd,cn,ads,cerdeno,bk,bbb,bno,bhn}. Here we evaluate the neutralino relic abundance $\\Omega_{\\chi} h^2$ and the neutralino-nucleon scalar cross-section $\\sigma_{\\rm scalar}^{(\\rm nucleon)}$, which is relevant to dark matter direct detection. In Sect.\\ref{sec:model} we define the supersymmetric scheme adopted in the present paper and in Sect. \\ref{sec:cross} we provide analytical considerations and numerical evaluations. Our conclusions are drawn in Sect. \\ref{sec:finale}. ",
        "conclusions": "\\label{sec:finale}  In the present paper we have focussed our attention on relic neutralinos of light masses: $m_{\\chi} \\lsim $ 45 GeV, which are allowed in supersymmetric models where no unification of gaugino masses is assumed. We have shown that these neutralinos may have elastic cross-sections off nucleons which go up to $\\sigma_{\\rm scalar}^{(\\rm nucleon)} \\sim 10^{-7}$ nbarn, with a relic abundance of cosmological interest: $0.05 \\lsim \\Omega_{\\chi} h^2 \\lsim 0.3$.  The present upper limits to $\\xi \\; \\sigma_{\\rm scalar}^{(\\rm nucleon)}$ provided by WIMP direct detection experiments \\cite{morales,dama0,cdms,edel} do not constrain the supersymmetric configurations for the light neutralinos considered here. This is especially true once the relevant uncertainties (mainly related to the form and parameters of the WIMP galactic distribution function \\cite{bcfs} and to the quenching factors for bolometric detectors) are taken into account. The CDMS upper bound \\cite{cdms} could concern a small fraction of supersymmetric configurations in the range around 15 GeV, though very marginally, if the uncertainties on astrophysical quantities are considered. Moreover, the CDMS bound needs a confirmation by a further running in a deep--underground site, as planned by the Collaboration.  The small--mass neutralino configurations analysed in the present paper are accessible to experiments of direct detection with a low--energy threshold and a high sensitivity.  An experiment of this type is the DAMA experiment with a mass of $\\simeq$ 100 kg of NaI(Tl), whose results after a 4-years running show an annual-modulation effect at a $4\\sigma$ C.L. which does not appear to be related to any possible source of systematics \\cite{damapl}. The DAMA experiment, with its high sensitivity, is potentially good to investigate also the relic neutralinos considered in the present paper."
    },
    "0212/astro-ph0212016_arXiv.txt": {
        "abstract": "There is growing evidence that soft gamma-ray repeaters (SGRs) and anomalous X-ray pulsars (AXPs) are isolated neutron stars with superstrong magnetic fields, i.e., magnetars, marking them a distinguished species from the conventional species of spindown-powered isolated neutron stars, i.e., radio pulsars. The current arguments in favor of the magnetar interpretation of SGR/AXP phenomenology will be outlined, and the two energy sources in magnetars, i.e. a magnetic dissipation energy and a spindown energy, will be reviewed. I will then discuss a missing link between magnetars and pulsars, i.e., lack of the observational evidence of the spindown-powered behaviors in known magnetars. Some recent theoretical efforts in studying such behaviors will be reviewed along with some predictions testable in the near future. ",
        "introduction": "For a long time, radio pulsars have been regarded as the only manifestation of isolated neutron stars\\footnote{The internal compositions and equations-of-state of ``neutron stars'' are not well determined. These stars could be in principle more exotic, e.g., could be composed of pure strange quark matter (e.g. Xu 2002, in these proceedings). Here I refer to ``neutron stars'' as a broader class of objects that includes more exotic categories.}. Recent observational developments indicate that isolated neutron stars also manifest themselves as other species (Pavlov et al. 2002, for a review), among which soft gamma-ray repeaters (SGRs) and anomalous X-ray repeaters (AXPs) have attracted growing attention in the neutron star community. These two types of objects originate, respectively, from the anomalous species of two distinct classes of phenomenon, i.e., gamma-ray bursts and accreting X-ray pulsars, but share many common features. Recently, two observational facts finally connect a bridge between SGRs and AXPs. First, after being quiescent for more than twenty years, SGR 0526-66 is found to have a steep non-thermal spectrum in the quiescent state which is similar to the non-bursting AXPs (Kulkarni et al. 2003). Second, soft, repeating bursts were recently detected from two AXPs, 1E 1048-5937 (Gavriil, Kaspi \\& Woods 2002) and 1E 2259+586 (Kaspi \\& Gavriil 2002). These suggest that SGRs/AXPs belong to a unified class of objects.  In the literature, there exist essentially four types of models to interpret SGR/AXP phenomenology. These are, according to the sequence of popularity, the magnetar model (Duncan \\& Thompson 1992; Paczy\\'nski 1992; Thompson \\& Duncan 1995, 1996; Thompson, Lyutikov \\& Kulkarni 2002), the accretion model involving fossil disks (Chatterjee et al. 2000; Alpar 2001; Masden et al. 2001), the models involving strange quark stars (Alcock et al. 1986; Cheng \\& Dai 1998; Zhang, Xu \\& Qiao 2000; Usov 2001), and the models involving magnetic white dwarfs (Paczy\\'nski 1990; Usov 1993). It is fair to say that at the current stage none of the models can interpret {\\em all} SGR/AXP observations satisfactorily. Nonetheless, the magnetar model has its merit to interpret most observations under one single hypothesis, i.e., SGRs/AXPs are neutron stars with superstrong magnetic fields ($\\sim 10^{14}-10^{15}$ G at the surface). Other models either have troubles to interpret some observations (e.g. the accretion model fails to account for the super-Eddington SGR bursts) or have to introduce additional assumptions to account for data (e.g. Zhang 2002a for a review). Below I will list the solid observational facts of SGR/AXPs and confront them with the magnetar model.  1. {\\bf Timing properties.} Known SGRs/AXPs exclusively have long periods [$P\\sim (5-12)$ s] and large spindown rates [$\\dot P \\sim 5\\times (10^{-13}-10^{-10})$ s/s]. Assuming magnetic braking, this directly refers to a superstrong surface magnetic fields [$B_s \\sim (10^{14}-10^{15})$ G if these objects are neutron stars. Irregular spindown may be a common feature of these objects, and is not necessarily related to the bursting behavior. This could be accomendated in a magnetar model with twisted magnetosphere (Thompson et al. 2002).  2. {\\bf Quiescent emission properties.} SGRs/AXPs all display a steady luminous X-ray emission with $L_x \\sim (10^{35}-10^{36})$ ergs/s, which could be explained in terms of magnetic dissipation (magnetic field decay, Thompson \\& Duncan 1996; or magnetic enhanced cooling, Heyl \\& Hernquist 1998; or untwisting of a global current-carrying magnetosphere, Thompson et al. 2002). Optical/IR counterparts have been detected from three AXPs (4U 0142+61, 1E 2259+586, and 1E 1048.1-5937), but no promising interpretation within the magnetar model is proposed. No gamma-ray and radio emission has been firmly detected from the SGRs/AXPs.  3. {\\bf Burst properties.} SGR bursts are soft and repeating, with luminosity ranging from $10^{38}$ ergs/s all the way up to $\\sim 10^{45}$ ergs/s (usually super-Eddington, and two most luminous bursts, namely giant flares, have been detected from SGR 0526-66 on March 5, 1979; and from SGR 1900+14 on August 27, 1998). A strength of the magnetar model is that it can interpret the bursting phenomenology successfully in terms of the magnetic cataclysmic dissipation events in superstrong magnetic fields. Super-Eddington bursts are natural in strong fields in which the Thomson cross section is suppressed.  4. {\\bf Environmental effects.} Most SGRs/AXPs are located close to supernova remnants (SNRs) in projection. Solid associations with the SNRs are yet firmly established. Real associations are consistent with the magnetar theory which predicts that these objects are young neutron stars, but the SNR ages are not fully consistent with the spindown age of these objects. Assuming associations, SGRs have larger proper motions than AXPs. That one AXP with SNR association, 1E 2259+586, recently displayed hundreds of repeating bursts make the issue more complicated. The claim that SGRs/AXPs are born in dense environments (Marsden et al. 2001) is not confirmed (Gaensler et al. 2001).  5. {\\bf Cyclotron features}. Cyclotron features have been detected in SGR outbursts (Ibrahim et al. 2002), which is consistent with the magnetar model if the features are of proton-origin, but refers to a much lower magnetic field if the features are of electron-origin.  In summary, though not fully unquestionable, the magnetar model is successful in many respects in interpreting the data. However, there is hitherto no definite proof that SGRs/AXPs are isolated neutron stars powered by superstrong magnetic fields. I believe that the key to prove the magnetar interpretation would be looking for a missing link between magnetars and normal pulsars, which I will lay out in the next section. ",
        "conclusions": ""
    },
    "0212/astro-ph0212220_arXiv.txt": {
        "abstract": "We present the detection and analysis of line index variations in the pulsating sdB star PG\\,1605+072. We have found a strong dependence of line index amplitude on Balmer line order, with high-order Balmer line amplitudes up to 10 times larger than H$\\beta$. Using a simple model, we have found that the line index may not only be dependent on temperature, as is usually assumed for oscillating stars, but also on surface gravity. This information will provide another set of observables that can be used for mode identification of sdBs. ",
        "introduction": "\\label{sec:intro-ew}  Subdwarf B (sdB) stars are hot, core-helium burning stars, with hydrogen envelopes that are too thin to sustain nuclear fusion. The discovery of multimode pulsations in some sdBs should allow the use of asteroseismology to probe their atmospheres, and thereby help to answer the many questions remaining about sdB formation and evolution.  Until recently, time-series observations of pulsating sdBs have been limited to photometry. We have shown that it is possible to detect radial velocity variations of Balmer lines in the pulsating sdB star PG\\,1605+072 using 2\\,m-class telescopes (\\citetwobare{OBK00b}{OBK02}, hereafter \\pIandII). Four-metre class telescopes have also been used to detect velocity variations in sdBs \\cite{JP00,WJP01}. In \\pII\\ we found closely spaced peaks in the amplitude spectrum of PG\\,1605+072, as well as amplitude variation in at least one of these peaks. In this paper we describe the analysis of line-index variations of Balmer lines in PG\\,1605+072. The atmospheric parameters of this sdB have been determined from high-resolution optical spectra and NLTE model atmospheres to be $T_\\mathrm{eff}$=32\\,300$\\pm$300\\,K and log\\,$g$=5.25$\\pm$0.05 \\cite{HRW99}. The helium abundance was found to be subsolar, i.e., log\\,(He/H)=$-$2.53$\\pm$0.1. The star was also found to be rotating with a projected rotational velocity of 39\\,km\\,s$^{-1}$.  The use of equivalent widths of Balmer lines to monitor stellar oscillations of solar-type stars was first proposed by \\citeone{KBV95}. The Balmer lines are very temperature sensitive -- the equivalent width of these lines depends on the number of hydrogen atoms in the second level. According to \\citeone{Kurucz79}, the maximum line strength of H$\\alpha$ to H$\\delta$ is between 8500\\,K and 9000\\,K at log\\,$g$=4.0. Since the temperature changes slightly with the oscillations, the equivalent width must also change. The idea was extended by \\citeone{BKR96} to include metal lines. In this paper we calculate \\textit{line indices}, which depend on the temperature and surface gravity sensitivity of the equivalent widths of spectral lines. This method has been used to identify modes in the $\\delta$ Scuti stars FG Vir \\cite{VKB98} and BN Cnc \\cite{DFL02}, and to detect equivalent-width oscillations in the H$\\alpha$ line of the roAp star $\\alpha$ Cir \\cite{BVB99}. Here we present the first detection of line index variations in a pulsating sdB star, namely PG\\,1605+072, and find a dramatic dependence of line-index amplitude on Balmer line number. We show that a simple model can qualitatively explain this effect as a combination of temperature and surface gravity fluctuations. ",
        "conclusions": "We have detected line index variations in the pulsating sdB star PG\\,1605+072. There is a strong dependence of line-index oscillation amplitude on Balmer line. We have developed a simple model, assuming hydrostatic equilibrium, where we use model spectra at as close to the same resolution and sampling of the observations as possible. Using this model, we have shown that despite not measuring true equivalent width variations, we can use the observables to infer effective temperature and effective surface gravity variations in PG\\,1605+072.  Using more detailed models, we should be able to measure the $\\Delta T$ and $\\Delta$log\\,$g$ independently of photometry and velocity measurements. With this information, we will have another set of amplitudes which we can use for mode identification of sdB stars.  \\vspace{0.3cm}  The authors would like to thank Mike Ireland for helpful discussions on error ellipses, and Steve Kawaler for continuing fruitful discussions on sdB evolution and interiors. This work was supported by an Australian Postgraduate Award (SJOT), the Australian Research Council, the Danish National Science Research Council through its Center for Ground-based Observational Astronomy, and the Danish National Research Foundation through its establishment of the Theoretical Astrophysics Center."
    },
    "0212/astro-ph0212199_arXiv.txt": {
        "abstract": "There are observational facts and theoretical arguments for an origin of gamma-ray bursts (GRBs) in molecular clouds in distant galaxies. If this is true,  one could detect a significant flux of GRB prompt and early afterglow X-ray radiation scattered into our line of sight by the molecular and atomic matter located within tens of parsecs of the GRB site long after the afterglow has faded away. The scattered flux directly measures the typical density of the GRB ambient medium. Furthemore, if the primary emission is beamed, the scattered X-ray flux will be slowly decreasing for several months to years before falling off rapidly. Therefore, it should be possible to estimate the collimation angle of a burst from the light curve of its X-ray echo and a measured value of the line-of-sight absorption column depth. It is shown that detection of such an echo is for the brightest GRBs just within the reach of the Chandra and XMM-Newton observatories. ",
        "introduction": "\\label{intro}  One of the fundamental questions relating to the gamma-ray burst (GRB) phenomenon is how much energy is released during a GRB. Given a measured burst fluence, the inferred energy release is proportional to $\\theta_0^2/2$, where $\\theta_0$ is the opening angle of the collimated relativistic fireball (Piran 1999) that is probably small ($\\theta_0\\ll 1$). Apart from constraining the burst energetics, $\\theta_0$ also determines the event rate of GRBs in the universe (which is proportional to $\\theta_0^{-2})$. Thus, the beaming angle is a very important quantity in GRB research.  There have been attempts to estimate $\\theta_0$ via observing a break in GRB afterglow light curves and interpreting this break as due to a slowing down of the jet to $\\Gamma<\\theta_0^{-1}$ (here $\\Gamma$ is the bulk Lorentz factor of the jet). Summarizing the information on jet break times in more than a dozen of GRBs Frail et al. (2001) have inferred values for $\\theta_0$ ranging from $1^\\circ$ to more than $25^\\circ$. These authors came to the conclusion that the gamma-ray energy release is narrowly clustered around $5\\times 10^{50}$~ergs. However, since their derivation was based on a particular interpretation of the observed breaks and that other explanations for such breaks exist (as mentioned by Frail et al. 2001), these measurements of $\\theta_0$ should necessarily be considered tentative for the time being.  Another fundamental issue is the environment in which GRBs occur. High gas column densities ($\\Nh\\sim 10^{22}$~cm$^{-2}$) toward GRB locations have been inferred from a spectral analysis of a sample of X-ray afterglows observed with BeppoSAX (Owens et al. 1998; Galama et al. 2001; Reichart \\& Price 2002). Such values of $\\Nh$ are typical of Galactic giant molecular clouds and therefore a strong indication that GRBs occur in star forming regions of their host galaxies, which in turn argues for collapse of massive stars as their sources (Woosley 1993; Paczy\\'{n}ski 1998).  In this paper we propose an observational test that enables direct determination (with some reservations) of both the collimation angle of a GRB and the typical density of the medium surrounding it on scales of pcs to tens of pcs typical of molecular clouds and complexes. This method consists of observing the location of a bright GRB with a powerful X-ray telescope several times during the first months to years after the burst. We predict that in such observations a weak X-ray flux may be detected, which will be radiation from the burst reaching us, delayed by scattering from the ambient medium.  The idea at the basis of the proposed method -- to search for scattered GRB radiation -- is not new. The theme of scattered X-ray emission was first introduced by Dermer, Hurley \\& Hartmann (1991) in the context of the then popular model of the galactic stellar binary origin of GRBs. Madau et al. (2000) considered Compton echoes from gamma-ray bursts arising in the circumburst environment within a fraction of a pc of the GRB site. These authors focused on the case where the primary burst emission is collimated away from us. Ramirez-Ruiz et al. (2001) furtherelaborated this scenario. Esin \\& Blandford (2000) studied the scattering of GRB early optical emission on the surrounding dust, while M\\'{e}sz\\'{a}ros \\& Gruzinov (2000) considered a similar phenomenon -- small-angle scattering of X-rays by dust.  Ghisellini et al. (1999) and B\\\"{o}ttcher et al. (1999) have computed the time dependence of fluorescence emission in the iron $K\\alpha$ line resulting from the interaction of the GRB radiation with the surrounding matter under the assumption that the burst emission is isotropic. Detection of such a fluorescent line would make it possible to determine the redshifts of GRBs with invisible optical afterglows. However, the predicted flux in the line is apparently too low to be detected with the current generation of X-ray telescopes unless the density of interstellar material at the GRB site is very high ($n\\gtrsim 10^4$~cm$^{-3}$). We note here (see \\S\\ref{effect}) that the scattered X-ray flux in the relevant 0.3--5~keV band will typically be an order of magnitude higher than the flux of fluorescence emission in the same band. As we shall see, this enhancement factor proves crucial for the problem at hand, as it makes the emergent signal detectable with the present-day X-ray telescopes provided that GRBs indeed originate within molecular clouds. ",
        "conclusions": "\\label{conclude}  We have described a new observational method that enables us to determine the typical ambient medium density of bright GRBs as well as their degree of collimation in an almost model-independent way. This method can hopefully be employed with the existing Chandra and XMM-Newton X-ray telescopes.  A suitable observational strategy would be to schedule an observation of the site of a very bright ($\\flu\\sim10^{-5}$~erg~cm$^{-2}$) GRB a few months after the burst. Should some X-ray flux be detected during this observation, an additional one or two observations should be carried out several months or years later so that the light curve of the scattered emission could be roughly reconstructed and the GRB opening angle could be estimated [from equation (\\ref{theta0})]. If the first observation fails to yield a significant flux, an interesting upper limit, $\\sim 10^3$~cm$^{-3}$, on the density of the medium surrounding the GRB on parsec scales can be derived [from equation (\\ref{dens})]. It will take a  $10^5$~s observation with Chandra or XMM-Newton to detect a scattered flux of a few $\\times 10^{-16}$~erg~cm$^{-2}$~s$^{-1}$ at 0.3--5~keV. We emphasize that it is sufficient to collect a total of several photons from the GRB direction, as no detailed spectral information is needed. We also note that a flux of $5\\times 10^{-16}$~erg~cm$^{-2}$~s$^{-1}$ corresponds to an equivalent isotropic luminosity of order $10^{42}$~erg~s$^{-1}$ for a burst at $z\\sim 1$. Therefore, the scattered GRB emission should outshine the X-ray emission of any non-active host galaxy.  Since GRBs with $\\flu\\sim10^{-5}$~erg~cm$^{-2}$ are very rare events (at best a few per year), constant monitoring of the whole sky with an instrument sensitive to X-rays is crucial. After the end of the CGRO mission, such a capability will be provided by the Swift satellite\\footnote{http://swift.gsfc.nasa.gov/}."
    },
    "0212/astro-ph0212150_arXiv.txt": {
        "abstract": "We present an overview of all the cosmic shear results obtained so far. We focus on the 2-point statistics only. Evidences supporting the cosmological origin of the measured signal are reviewed, and issues related to various systematics are discussed. ",
        "introduction": "The cosmic shear is a gravitational lensing effect caused by the large scale structures in the universe on the distant galaxies. Its net effect is to distort coherently the galaxy images over large angular scales and to change the focusing properties of the light beam (galaxies appear larger or smaller depending on the intervening mass density). Only the former effect is 'easily' measurable, and is the one we will discuss here. The latter, which has been measured in some cluster lensing cases, is much more challenging for cosmic shear purposes.  The measurement of the amplitude of the cosmic shear as a function of scale is a direct indicator of the projected mass power spectrum, convolved with a selection function which only depends on the cosmological parameters and the redshift distribution of the sources. It is therefore a tool for evaluating the mass distribution in the universe, as well as for measuring the cosmological parameters. In the following, we first outline the theory of cosmic shear, and its link to measurable quantities in Section2. We discuss the measurements and the systematics in Section 3. Section 4 is a discussion of the cosmological parameters measurements and the still open problems associated with it. ",
        "conclusions": "As quoted earlier \\cite{j97}, the cosmic shear signal depends primarily on four parameters: the cosmological mean density $\\Omega_m$, the mass power spectrum normalization $\\sigma_8$, the shape of the power spectrum $\\Gamma$, and the redshift of the sources $z_s$. Therefore, any measurement of the statistics shown in Section 2 provide constraints on these parameters. A big enough lensing survey \\cite{huteg99} will also provide constraints on many other parameters, but we do not discuss this issue here.  \\begin{figure}[t] \\centerline{ \\psfig{figure=flat_omega_sigma8.ps,height=6.cm} \\psfig{figure=contour_OM_S8_flat.eps,height=6.cm} } \\caption[]{The solid lines on each plot show the 1, 2 and 3$\\sigma$ contours of the VIRMOS and RCS survey, from the measurements shown in Figure \\ref{mapplot}. The contours have been marginalised over the source redshift and the slope of the matter power spectrum as described elsewhere \\cite{vwal01,hoekstra02}. \\label{flat_omega_sigma8.ps}} \\end{figure} Figure \\ref{flat_omega_sigma8.ps} shows the joint $\\Omega_m$, $\\sigma_8$ constraints obtained from the measurements of Figure \\ref{mapplot}. They are obtained only when comparing the measured lensing signal to the non-linear predictions. Indeed, the actual surveys are not yet big enough to probe the linear scales accurately. The non-linear power can be computed numerically \\cite{smith02}, but its precision is still uncertain. Recent investigations show that a $10 \\%$ r.m.s. uncertainty is expected, which means that the cosmological parameters cannot be known with better precision for the moment. According to the Figures \\ref{model.ps}, \\ref{kernel.ps}, the transition scale between the linear and non-linear regimes is around 1 degree. The consequence is that the quoted mass normalization $\\sigma_8$ is sensitive to the validity of the non-linear mapping at small scale. In this respect, \\cite{jar02} are less contaminated by this problem because they used the lensing signal from $30'$ to $100'$ to constrain the mass normalization.  Table \\ref{constraints} summarizes the $\\sigma_8$ measurements for all the lensing surveys published so far. For simplicity it is given for $\\Omega_m=0.3$. Despite the differences among the surveys, it is worth to note that the results are all consistent within $2.5\\sigma$ between the most extreme cases, when poorly known parameters are marginalised.  The residual $B$ mode has been included in the analysis of the most recent cosmic shear surveys \\cite{vwal02, hoekstra02,ham02,jar02}, but we do not know yet how to deal with it properly: some groups \\cite{vwal02,hoekstra02,ham02} added the $B$ mode in quadrature to the $E$ errors, taking into account the correlation between various scales. The $B$ mode has been substracted first from the $E$ mode in \\cite{hoekstra02}, and not in \\cite{vwal02}. This might probably result in a slight bias for high $\\sigma_8$ values in \\cite{vwal02}. Unfortunately we have no garantee that the $B$ substraction is the right correction method. Recently \\cite{jar02} marginalised the probabilities over $E-B$ to $E+B$ taken as the signal, which is more likely to include the 'true' $B$ mode correction one has to apply.  The $B$ mode correction and the non-linear power spectrum predictions are, today, the major limitations of cosmic shear surveys. They prevent the measurements of cosmological parameters below $10\\%$, even if today's lensing surveys could do better in principle. It is therefore very difficult to establish a discrepancy among the cosmic shear results today. There is probably something to learn about basic PSF correction with the existing data, since the different groups get different $B$ mode amplitude and scale dependence. This is still widely unexplored.  The positive aspect is that cosmic shear is now established as a powerfull tool to probe the dark matter and the cosmological parameters, and that it works, which was not the case only 2 years ago. There is a lot to expect from forthcoming surveys, with the first results expected in less than a year from now."
    },
    "0212/astro-ph0212366_arXiv.txt": {
        "abstract": "The oldest stars hide information about the chemical evolution of the early universe. In this contect new models of the AGB stellar evolution phase at very low and zero metallicity are presented. Due to the deficiency or absence of CNO catalytic material hydrogen burning operates at significantly higher temperatures than in stars of higher metallicity. As a result convective mixing plays a very important role since the nuclear burning shells are not well separated by an entropy barrier. These stars can host flash-like burning events induced by mixing of $\\czw$ and protons at the core-envelope interface. Three different mixing events with proton-capture nucleosynthesis can be encountered. One of them, the hot dredge-up, is reported here for the first time. All these specific modes of nucleosynthesis are relevant for the production of CNO material in the first intermediate mass stars. ",
        "introduction": "The first nuclear processing of pristine Big Bang matter has left a signature in today's most metal poor low mass stars. There is some debate whether both low mass and massive stars contributed to this processing (Nakamura \\& Umemura, 2002) or whether the first stars were exclusively very massive (e.g.\\ Abel et al., 2002). The recently discovered low mass star \\mbox{HE0107-5240} has [Fe/H]=-5.3 (Christlieb et al., 2002). This may indicate that the  Pop.\\,III IMF has a low mass component. The large nitrogen overabundance in this object could mean that intermediate mass stars are involved in the contamination of this star.  This is much more obvious for stars like \\mbox{HE0024-2523} (Lucatello et al., 2003), a very metal poor main sequence star with [Fe/H]=-2.7, an \\spr\\ abundance signature, large C and N overabundances as well as radial velocity variations. It has been suggested that this star and related objects have been polluted by an intermediate mass binary companion. The carbon and \\spr\\ abundance could -- in principle -- be understood with the dredge-up and \\spr\\ mechanisms known from AGB stars of solar-like metallicity (e.g.\\ Gallino et al., 1998). The nitrogen overabundance  could result from partial envelope CNO cycling (hot bottom burning, HBB). Siess et al.\\ (2002) predict that both carbon as well as nitrogen may be significantly enhanced due to dredge-up and HBB in Z=0 AGB stars. Their $3(2)\\msun$ model predicts at the last thermal pulse (TP)  [C/H]$\\sim -0.9(0.0)$, [N/H]$\\sim +0.4(-2.2)$ and C/O$=5(53)$. \\mbox{HE0024-2523} has  [C/H]$= -0.1$, [N/H]$= -0.6$ and C/O$=100$. A [C/N] ratio of a few tenth of a dex is also found in other similar stars (e.g.\\ \\mbox{CS 22948}, Aoki et al., 2002). It is even more extreme in \\mbox{HE0107-5240}. This puts some limit on the role HBB can play because efficient HBB may too quickly decrease the C/N ratio below unity. \\begin{figure} \\plotone{FF08348.ps} \\caption{\\label{fig:FF08348} Proton ingestion and burning phase during the 10$^{\\rm th}$ TP of a Z=0 sequence with M=$5\\msun$. Top panel: Abundance profile of upper He-flash convection zone (left), H-flash convection zone (middle) induced by proton ingestion from above, and H-shell ashes on top of the core (right). Bottom panel: Convective mixing efficiency in terms of a diffusion coefficient.} \\end{figure}   The role of intermediate mass stars as a source of primary nitrogen in the early chemical evolution has recently been revisited on the basis of the N abundances in Damped Lyman-$\\alpha$ Absorbers (Pettini et al., 2002, and Henry \\& Proschaska, this volume). Meynet \\& Maeder (2002) have shown that in intermediate mass stars of very low metallicity rotationally induced mixing can lead to significant $\\nvi$ enrichment of the outer layers even before the first TP occurs. In these models the C/N ratio is much smaller than unity and would not reproduce the abundance patterns of stars like \\mbox{HE0024-2523}. Clearly, the  the envelope abundances are modified in the subsequent TP AGB phase.  In  very low and zero metallicity stars the convection zone driven by the He-flash may reach into the proton rich unprocessed outer layers . As protons enter the He-shell environment they initiate a vigorous non-equilibrium nuclear burning, mainly of proton captures on $\\czw$. This may happen in the case of the He-core flash as well as the AGB He-shell flash for low enough metallicity and core mass (Fujimoto et al., 2000). The studies by Siess et al.\\ (2002) and Chieffi et al. (2001) found that the TP AGB evolution of Z=0 objects eventually resembles  that of only mildly metal poor AGB stars. Either the second dredge-up or a convective instability of the H-shell at reignition after a TP sufficiently increases the catalyst material for the CNO cycle.  In this paper I present first results on stellar evolution calculations that consider convective nucleosynthesis in a numerically consistent way and also allow for soft convective boundaries with respect to mixing. Models of the proton ingestion induced H-flash during TPs of Z=0 AGB stars of $2\\msun$ and $5\\msun$ are described. In addition I report on a new process of carbon and nitrogen enrichment uniquely present in very low metallicity stars which is a combination of HBB and third dredge-up. \\begin{figure} \\plotone{MH-LH-M5Z0.0001.ps} \\caption{\\label{fig:MH-LH-M5Z0.0001} Evolution of the mass coordinate of the H-free core $M_{\\rm H}$ and the H-burning luminosity $L_{\\rm H}$ during the end of the second dredge-up ($t<7500{\\rm yr}$) and the first two AGB TPs of the Z=0.0001 sequence with M=$5\\msun$. }\\end{figure} ",
        "conclusions": "All recent calculations of the  AGB phase at Z=0 agree that thermal pulses and dredge-up occur. Various processes, in particular rotationally induced mixing during the He-core burning phase, can lift the envelope abundance of catalytic material for the CNO cycle so that eventually the evolution in these objects resemble those of only mild metal deficiency. It is therefore not clear if the proton ingestion into the He-shell flash convection zone or the HCE take place in real Z=0 AGB stars should they have existed. Here a hot variant of the third dredge-up was described which may be an important ingredient in understanding both the primary production of nitrogen as well as the peculiar carbon and nitrogen overabundances observed in very metal poor stars.  \\paragraph{Acknowledgments:} I would like to thank D.A. VandenBerg for support through his Operating Grant from the Natural Science and Engineering Research Council of Canada."
    },
    "0212/nucl-th0212072_arXiv.txt": {
        "abstract": "We study in a fully selfconsistent approach the structure of a vortex in low density superfluid neutron matter. We determine that the matter density profile of a vortex shows a significant depletion in the region of the core, a feature never reported for a vortex state in a Fermi superfluid. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/astro-ph0212400_arXiv.txt": {
        "abstract": "We consider the physical conditions under which supermassive black holes could have formed inside the first galaxies. Our SPH simulations indicate that metal-free galaxies with a virial temperature $\\sim 10^4$ K and with suppressed H$_2$ formation (due to an intergalactic UV background) tend to form a binary black hole system which contains a substantial fraction ($\\ga 10\\%$) of the total baryonic mass of the host galaxy. Fragmentation into stars is suppressed without substantial H$_2$ cooling.  Our simulations follow the condensation of $\\sim 5\\times 10^6 M_\\odot$ around the two centers of the binary down to a scale of $\\la 0.1$~pc.  Low-spin galaxies form a single black hole instead.  These early black holes lead to quasar activity before the epoch of reionization.  Primordial black hole binaries lead to gravitational radiation emission at redshifts $z\\ga 10$ that would be detectable by LISA. ",
        "introduction": "Supermassive black holes (SMBHs) are believed to provide the power source of quasars via the accretion of surrounding gas (Salpeter 1964; Zeldovich 1964; Lynden-Bell 1969; Rees 1984).  They were able to form already at early cosmic times, as implied by the recent discovery of quasars at redshifs $z\\ga 6$ (Becker et al. 2001; Djorgovski et al. 2001; Fan et al. 2002).  The existence of SMBHs with inferred masses of $\\ga 10^{9}M_{\\odot}$, less than a billion years after the big bang, provides important constraints on any formation scenario (Haiman \\& Loeb 2001).  The gas physics involved in the formation of SMBHs is still not well understood (Rees 1984, 2002; Loeb \\& Rasio 1994; Barkana \\& Loeb 2001).  In this work, we attempt to simulate this process under the simplest and best prescribed set of initial conditions; those dictated by the early universe.  The earliest galaxies are simple in that they are the first condensations of gas to grow out of the seed inhomogeneities in the early universe.  The composition of the primordial gas is determined by big bang nucleosynthesis (Burles, Nollett, \\& Turner 2001, and references therein) and any primordial magnetic fields are not expected to be dynamically significant.  Previous numerical simulations of early luminous structures (Bromm, Coppi, \\& Larson 1999, 2002; Abel, Bryan, \\& Norman 2000, 2002; Nakamura \\& Umemura 2001) have focused on the formation of stars inside the very first gaseous objects with a mass, $\\sim 10^5M_\\odot$, just above the cosmological Jeans mass.  Due to their low virial temperature (hundreds of K), fragmentation into stars is possible inside these objects through the formation of molecular hydrogen, H$_2$, which cools efficiently via rotational-vibrational transitions even at these low temperatures. However, a relatively modest UV flux is sufficient to photo-dissociate the fragile H$_2$ molecules and suppress their role in cooling the gas (Haiman, Rees, \\& Loeb 1997; Ciardi et al. 2000; Haiman, Abel, \\& Rees 2000).  A destructive flux of UV photons could be produced by a small, early population of stars.  Here we consider dwarf galaxies with virial temperatures $\\ga 10^4$~K inside of which atomic cooling is effective (Oh \\& Haiman 2002). These galaxies comprise a substantial population at $z\\sim 10$ (Barkana \\& Loeb 2001) that is expected to dominate the reionization of hydrogen at $z\\ga 6$ (Wyithe \\& Loeb 2003a).  We use smoothed particle hydrodynamics (SPH) simulations to describe the cooling and dynamics of the gas for different choices of its total angular momentum (or spin parameter), following the suggestion by Eisenstein \\& Loeb (1995) that low-spin systems would be more susceptible to the formation of a SMBH.  If H$_2$ cooling is suppressed inside these galaxies, then their gas will not cool below $10^4$K. When the temperature to which the gas cools is only somewhat lower than the virial temperature of the host galaxy (at which the pressure force balances gravity for the gas), we expect that fragmentation into small clumps will be avoided and the gas will tend to condense isothermally (at its temperature floor of $\\sim 10^4$K) into large clumps. Such large clumps may then collapse to form a SMBH, possibly through an intermediate stage of a supermassive star.  The viability of this scenario relies on the suppression of molecular H$_2$ cooling, which when present is capable of cooling the gas to a temperature as low as 200 K.  We therefore start our analysis in \\S~2 by considering the cosmic UV background expected from the first stars. Later on, in \\S~5 we carry out numerical simulations to determine the level of UV radiation necessary for the photodissociation of H$_2$ molecules during the early collapse phase of the above dwarf galaxies.  The SPH code we use is the same as was used to simulate the formation of the first massive stars (Bromm et al. 1999, 2002), except that we apply it to more massive galaxies where atomic cooling is effective. Our goal is to find whether there is a direct route to the formation of SMBHs under these circumstances.  Although our focus is on gas dynamics, we note that the first massive black holes (BHs) may have also formed out of the gravitational dynamics of clusters of stellar-mass BHs (Larson 2000, 2002; Madau \\& Rees 2001; Schneider et al. 2002; Islam, Taylor, \\& Silk 2003; Volonteri, Haardt, \\& Madau 2003).  Throughout this paper, we assume a standard $\\Lambda$CDM cosmology, with a total density parameter in matter of $\\Omega_{m}=1- \\Omega_{\\Lambda}=0.3$, and in baryons of $\\Omega_{\\rm B}=0.045$.  The Hubble constant is $h=H_{0}/(100$ km s$^{-1}$ Mpc$^{-1}$)=0.7, and the present-day power-spectrum amplitude is $\\sigma_{8}=0.9$ in spheres of radius $8h^{-1}$Mpc. ",
        "conclusions": "Our numerical simulations show that metal-free dwarf galaxies at redshifts $z\\sim 10$ whose cooling is dominated by atomic transitions (with a virial temperature just above $10^4$K) tend to form a central condensation consisting of one or two clumps which contain $\\ga 10\\%$ of the total baryonic mass of the galaxy.  As long as H$_{2}$ formation is suppressed, these massive clumps do not fragment but rather cool and continue to collapse isothermally at a temperature of $\\sim 10^4$K.  We have simulated the collapse to the point where $\\ga 10^6M_\\odot$ of gas condense to a scale $\\la 1$pc. At the end of our simulation the clump maintains a nearly free-fall collapse and shows no signs of fragmentation. We expect the collapse to continue until the gas cloud becomes optically thick to Thomson scattering and its cooling time is much longer than the infall time. At this stage the collapse may be halted due to viscous dissipation of rotational energy, and a rotating SMS is likely to form. After shedding a small fraction of its mass, the SMS will collapse to a SMBH (Baumgarte \\& Shapiro 1999; Shibata \\& Shapiro 2002).  The possible existence of a supermassive star as an intermediate stage in the formation of the first quasars can be tested by direct observations.  The SMS spectrum is expected to be close to thermal at an effective temperature of $\\sim 10^5$~K (Bromm et al.  2001b). The SMS mass can be inferred from its bolometric luminosity, given that the latter equals the Eddington value, $L_{\\rm Edd}=1.4\\times 10^{44}~{\\rm erg~s^{-1}}(M/10^6M_\\odot)$.  The emission redshift can be easily derived from the Gunn-Peterson trough produced by Ly$\\alpha$ absorption of the neutral IGM (Gunn \\& Peterson 1965).  Our model predicts that the emission spectrum of the SMS would show only hydrogen and helium lines but no metal lines (such as the broad emission lines observed for quasars at lower redshifts), since metal enrichment would signal small scale fragmentation into stars and would in turn enable the production of molecules on dust grains.  However, all the above characteristics are common to clusters of Population~III stars (Bromm et al. 2001b). To distinguish between a SMS and the extended image of a star cluster requires high angular resolution, possibly achievable with the {\\it James Webb Space Telescope}\\footnote{See http://ngst.gsfc.nasa.gov/}. Alternatively, it is possible to make use of the high probability of gravitational lensing by intervening galaxies out to $z\\sim 10$ (Barkana \\& Loeb 2000; Wyithe \\& Loeb 2002a,b; Comerford, Haiman, \\& Schaye 2002).  In particular, gravitational microlensing by stars within the lens galaxies (Wyithe \\& Loeb 2002b) is only possible for a SMS and not for a star cluster whose extent is much larger than the Einstein radius of a solar-mass star at a cosmological distance, $\\sim 10^{-2}~{\\rm pc}$.  Our calculations show that fragmentation into stars is suppressed inside the above dwarf galaxies only if H$_2$ cooling is negligible. The photodissociation of H$_2$ molecules requires a background flux of UV photons below 13.6eV (to which the neutral IGM is nearly transparent) that could possibly be produced by stars at $z\\sim 10$, prior to reionization (see \\S 2).  {\\it What is the comoving density of the first generation of massive BHs?}  At $z\\sim 10$, the comoving density of halos with a virial temperature $\\ga 10^4$K is $\\sim 5~{\\rm Mpc}^{-3}$ (Sheth \\& Tormen 1999).  Requiring that the total mass density of massive BHs would not exceed the value observed in the local universe (Yu \\& Tremaine 2002), $\\sim 2.5\\times 10^5M_\\odot~{\\rm Mpc^{-3}}$, we infer that the average BH mass per halo must be $\\la 5\\times 10^{4}M_\\odot$, much smaller than the characteristic mass of the gas clumps in our simulations. This follows naturally from our expectation that only a small fraction of all halos are both metal poor and exposed to a high level of UV flux during their entire history (see \\S 2 and \\S 5).  Even though a substantial fraction of the IGM volume is not enriched with metals at $z\\sim 10$ (Thacker et al. 2002; Furlanetto \\& Loeb 2003), most $2\\sigma$ halos may have been enriched by stars that were formed inside their hierarchical building blocks at an earlier cosmic time, when the UV background was much lower. The metals may then facilitate H$_2$ formation on dust particles as well as efficient cooling through metal emission lines (both atomic and molecular), resulting in further fragmentation.  In such a case, only a small minority of all halos forms by direct accretion of pristine gas from the IGM, as we have assumed in our simulations.  Alternatively, the final BH may contain a small fraction of the initial mass of its parent clump of gas.  Since we did not follow the collapse of each clump down to its Schwarzschild radius, we cannot be certain of the final BH formation efficiency.  It is possible that hydrodynamic or radiative feedback from a growing BH seed at the center limits the final mass that it obtains, so that eventually most of the mass of its parent clump is expelled in a wind due to energy release near the center (e.g., Haehnelt, Natarajan, \\& Rees 1998; Silk \\& Rees 1998).  The correlation between BH mass and the velocity dispersion of its host galaxy as observed in the local universe (Merritt \\& Ferrarese 2001; Tremaine et al. 2002) would imply an exceedingly low efficiency ($\\la 10^{-4}$) for the BH formation process in the dwarf galaxies under consideration (Wyithe \\& Loeb 2002c). However, this phenomenological correlation is measured in metal-rich galaxies where molecular cooling and star formation are abundant, and it may not hold under the unusual physical conditions inside the first generation of metal-free galaxies where the seeds for the present-day population of BHs were produced.  An early formation of SMBHs would lead to a considerable rate of detectable bursts of gravitational radiation from the coalescence of BH binaries at high redshifts (Wyithe \\& Loeb 2003b). In addition to the commonly discussed origin of binaries in galaxy mergers (a possible example of which has recently been observed in the nearby galaxy NGC~6240 by Komossa et al. 2003), we have demonstrated that BHs may form in binaries to start with, just as stars do."
    },
    "0212/astro-ph0212546_arXiv.txt": {
        "abstract": "In order to examine whether or not high-ionization nuclear emission-line regions (HINERs) in narrow-line regions of active galactic nuclei are dusty, we focus on two high-ionization forbidden emission lines, [Fe {\\sc vii}]$\\lambda$6087 and [Ne {\\sc v}]$\\lambda$3426. We perform photoionization model calculations to investigate possible dependences of the flux ratio of [Fe {\\sc vii}]$\\lambda$6087/[Ne {\\sc v}]$\\lambda$3426 on various gas properties, in order to investigate how useful this flux ratio to explore the dust abundances in HINERs. Based on our photoionization model calculations, we show that the observed range of the flux ratio of [Fe {\\sc vii}]$\\lambda$6087/[Ne {\\sc v}]$\\lambda$3426 is consistent with the dust-free models while that is hard to be explained by the dusty models. This suggests that iron is not depleted at HINERs, which implies that the HINERs are not dusty. This results is consistent with the idea that the HINERs are located closer than the dust-sublimation radius (i.e., inner radius of dusty tori) and thus can be hidden by dusty tori when seen from a edge-on view toward the tori, which has been also suggested by the AGN-type dependence of the visibility of high-ionization emission lines. ",
        "introduction": "The narrow-line region (NLR) is one of the fundamental ingredients of active galactic nuclei (AGNs) such as Seyfert galaxies, and thus its physical and chemical properties have been studied intensively up to now (see Osterbrock \\& Mathews 1986 for a review). Since the presence of dust grains significantly affects emission-line spectra in several ways, it has been often discussed whether or not dust grains survive in the NLRs. The chemical composition of the gas phase is modified by the depletion of refractory elements into the grains. Electrons photoelectrically ejected from the grains heat the gas, while electron-grain collisions cool it. The dust opacity modifies the transfer of the ionizing continuum that irradiates gas clouds in NLRs. Therefore the knowledge of dust abundances is required to interpret observed emission-line spectra of NLRs (see, e.g., Ferguson et al. 1997a).  There are some pieces of evidence that suggest the presence of dust grains in NLRs. Asymmetric narrow emission-line profiles may suggest the presence of dust grains in the gas (e.g., Heckman et al. 1981; De Robertis \\& Osterbrock 1984; Whittle 1985a, 1985b). In partially-ionized regions in NLRs the depletion of some refractory elements is suggested by the weakness of the [Ca {\\sc ii}] emission (e.g., Kingdon, Ferland, \\& Feibelman 1995; Ferguson et al. 1997a) and by comparing the intensity of near-infrared [Fe {\\sc ii}] emission lines with that of [O {\\sc i}]$\\lambda$6300 (e.g., Mouri et al. 1989; Simpson et al. 1996; Alonso-Herrero et al. 1997; Mouri, Kawara, \\& Taniguchi 2000) and with that of [P {\\sc ii}]1.188$\\mu$m (e.g., Oliva et al. 2001; Rodr\\'{\\i}guez-Ardila et al. 2002). Some resonance lines have been observed to be very weak, which seems to be also due to line transfer effect within dusty gas (e.g., Kraemer et al. 2000). These naturally lead to the following question; is the NLR ubiquitously dusty?  The innermost region of NLRs, which is irradiated by high photoionizing flux, is thought to be rather high-velocity, highly-ionized, dense gas clouds (e.g., De Robertis \\& Osterbrock 1986; Murayama \\& Taniguchi 1998a, 1998b; Barth et al. 1999; Tran, Cohen, \\& Villar-Martin 2000; Nagao, Taniguchi, \\& Murayama 2000b; Nagao, Murayama, \\& Taniguchi 2001a, 2001b). Since the gas clouds in this region can radiate very high ionization emission lines such as [Ne {\\sc v}], [Fe {\\sc vii}], and [Fe {\\sc x}] (the so-called coronal lines), such regions are called ``coronal line regions'' or ``high-ionization nuclear emission-line regions'' (HINERs: Binette 1985; Murayama, Taniguchi, \\& Iwasawa 1998). Nussbaumer \\& Osterbrock (1970) reported that the gas-phase abundance of iron that is estimated by using the emission-line flux ratio of [Fe {\\sc vii}]$\\lambda$6087/[Ne {\\sc v}]$\\lambda$3426 is close to the solar value. This means that the iron is not depleted significantly because of the absence of grains in HINERs. However, Ferguson, Korista, \\& Ferland (1997b) mentioned that the flux ratio of [Fe {\\sc vii}]$\\lambda$6087/[Ne {\\sc v}]$\\lambda$3426 depends on physical properties of gas clouds and thus it may be inappropriate to use this ratio as an indicator of the gas-phase iron abundance. They proposed alternative indicators of the gas-phase abundances of refractory elements by using some near-infrared high-ionization lines such as [Ca {\\sc viii}]2.32$\\mu$m, and then suggested that the HINERs seem to be dust-free (see also Ferguson et al. 1997a). Although their method is a robust one, it can be applied only to very few objects since the near-infrared high-ionization lines are hard to be measured. Indeed they applied their method only on the two Seyfert galaxies, NGC 1068 and the Circinus galaxy, and they concluded that the HINERs in the two Seyfert galaxies are dust free.  \\begin{figure*} \\epsscale{0.80} \\plotone{Nagao.fig1.eps} \\caption{ Frequency distributions of the two emission-line flux ratios, [Ne {\\sc v}]$\\lambda$3426/[O {\\sc iii}]$\\lambda$5007 and [Fe {\\sc vii}]$\\lambda$6087/[O {\\sc iii}]$\\lambda$5007, for the type 1 AGNs and the type 2 AGNs in our sample. \\label{fig1}} \\end{figure*}  In order to investigate whether the HINERs contain dust grains for a large sample of AGNs, we focus on [Fe {\\sc vii}]$\\lambda$6087 and [Ne {\\sc v}]$\\lambda$3426 again. These two forbidden emission lines are thought to arise at similar regions since their critical densities are similar ($1.6 \\times 10^7$ cm$^{-3}$ and $3.6 \\times 10^7$ cm$^{-3}$, respectively), and the ionization potentials of Fe$^{6+}$ and Ne$^{4+}$ are also nearly the same (97.1 eV and 99.1 eV, respectively). In addition to this advantage, it should be also an great advantage that these two emission lines are enough strong to be measured easily. Therefore, they are useful to examine the dust abundances in the HINERs if we know some possible dependences of their emissivities on physical properties, even though they might be inappropriate to determine exact gas-phase elemental abundances as noted by Ferguson et al. (1997b). In this paper, we investigate the dependences of the two high-ionization emission lines on physical properties based on photoionization model calculations. We then examine how these emission lines can give constraints on the issue whether or not the HINERs are dusty. ",
        "conclusions": "In this paper, we have compared the observed data on two high-ionization forbidden emission lines, [Fe {\\sc vii}]$\\lambda$6087 and [Ne {\\sc v}]$\\lambda$3426, with the results of photoionization model calculations, and found that the HINERs in NLRs of AGNs are not dusty. This finding was already noted by, e.g., Nussbaumer \\& Osterbrock (1970) and Ferguson et al. (1997b). However, our finding is important because we have shown the absence of dust grains in HINERs for a large sample of AGNs, for the first time.  In order to examine this conclusion more directly, we should carry out spatially resolved imaging observation of the thermal emission of dust in NLRs of AGNs. By using ALMA (Atacama Large Millimeter/submillimeter Array), we can resolve structures in the 0.01 arcsec scale by submillimetric imaging observation. This corresponds to $\\sim$0.02 pc for the Circinus galaxy (adopting the distance of 3.8 Mpc; Freeman et al. 1977). Thanks to this ability of high spatial resolution, we will be able to investigate the distribution of dust in NLRs directly."
    },
    "0212/astro-ph0212293_arXiv.txt": {
        "abstract": "X-ray emitting clusters of galaxies are considered in the context of modified Newtonian dynamics (MOND).  I show that self-gravitating isothermal gas spheres are not good representations of rich clusters with respect to the radial gas density distribution as indicated by the X-ray surface brightness. Pure gas spheres with a density distribution described by a $\\beta$ model, as observed, also fail because, with MOND, these objects are far from isothermal and have a gas mass and luminosity much larger than observed for clusters of the same mean temperature.  These problems may be resolved by adding an additional dark mass component in the central regions;  a constant density sphere contained within two core radii and having a total mass of one to two times that in the gas.  With this additional component, the observed luminosity-temperature relation for clusters of galaxies is reproduced. When the observed X-ray surface brightness distribution in actual clusters is modeled by such a two-component structure, the typical mass discrepancy is three to four times smaller than with Newtonian dynamics.  Thus while MOND significantly reduces the mass of the dark component in clusters it does not remove it completely. I speculate on the nature of the dark component and argue that this is not a fundamental problem for MOND. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/astro-ph0212552_arXiv.txt": {
        "abstract": "We report on the linear matter power spectrum reconstructed from the Ly$\\, \\alpha$ forest, on baryonic wiggles in the galaxy power spectrum, and on how to measure real space (as opposed to redshift space) power. \\vspace{1pc} ",
        "introduction": "Much of what I presented at the DARK2002 meeting is contained in \\cite{GH02,HT02,THX02,TZH01,Wang01}. To avoid redundancy, I will not repeat what is in those papers, but instead will devote this space to a subset of issues addressed at the meeting, namely the linear matter power spectrum reconstructed from the Ly$\\, \\alpha$ forest, baryonic wiggles in the galaxy power spectrum, and how to measure real space (as opposed to redshift space) power.  \\xiklyaconcordfig ",
        "conclusions": "Cosmologists now have a Standard Model, a flat $\\Lambda$CDM model seasoned with baryons. At the DARK2002 conference several speakers reported how data from diverse sources are pointing at the same concordance model.  Figure~\\ref{xiklyaconcord} illustrates one example of this concordance, how the linear power spectrum inferred from the Ly$\\, \\alpha$ forest is in full agreement with the concordance model.  First results from the new generation of large redshift surveys are emerging. It is probably premature to claim a detection of baryonic wiggles, but there are tantalizing hints of such, Figure~\\ref{xikbl}, and it should not be long before baryonic wiggles in the galaxy power spectrum become an observational reality.  This commentary concluded by bringing attention to the counter-intuitive fact that it is possible to measure the real space power spectrum of galaxies more accurately than the redshift space power spectrum. One hopes that those who analyse redshift surveys will begin to take advantage of this have-your-cake-and-eat-it truth."
    },
    "0212/astro-ph0212078_arXiv.txt": {
        "abstract": "We investigate the stellar disk properties of a sample of 19 nearby spiral galaxies with low inclination and late Hubble type (Scd or later).  We combine our high-resolution {\\it Hubble Space Telescope} I-band observations with existing ground-based optical images to obtain surface brightness profiles that cover a high dynamic range of galactic radius. Most of these galaxies contain a nuclear star cluster, as discussed in a separate paper.  The main goal of the present work is to constrain the properties of stellar bulges at these extremely late Hubble types.  We find that the surface brightness profiles of the latest-type spirals are complex, with a wide range in shapes.  We have sorted our sample in a sequence, starting with ``pure'' disk galaxies (approximately 30\\% of the sample).  These galaxies have exponential stellar disks that extend inwards to within a few tens of pc from the nucleus, where the light from the nuclear cluster starts to dominate. They appear to be truly bulge-less systems. Progressing along the sequence, the galaxies show increasingly prominent deviations from a simple exponential disk model on kpc scales. Traditionally, such deviations have prompted ``bulge-disk'' decompositions. Indeed, the surface brightness profiles of these galaxies are generally well fit by adding a second (exponential) bulge component. However, we find that most surface brightness profiles can be fit equally well (or better) with a single S\\'ersic-type $R^{1/n}$ profile over the entire radial range of the galaxy, without requiring a separate ``bulge'' component. We warn in a general sense against identification of bulges solely on the basis of single-band surface brightness profiles. Despite the narrow range of Hubble types in our sample, the surface brightness profiles are far from uniform. The differences between the various galaxies appear unrelated to their Hubble-types, thus questioning the usefulness of the Hubble-sequence for the sub-categorization of the latest-type spirals. A number of galaxies show central excess emission on spatial scales of a few hundred parsec that can be attributed neither to the nuclear cluster, nor to the S\\'ersic-type description of the stellar disk, nor to what one would generally consider to be a bulge component. The origin of this light component remains unclear. ",
        "introduction": "The question whether the morphology of galaxies is imprinted by the initial conditions of their formation or rather determined by secular evolution remains a subject of intense debate. The existence of the Hubble sequence has for many years provided important constraints on this issue. In very simple terms the Hubble sequence tells us that galaxies are made up of two components: a bulge and a disk. The canonical view of these components has long been that bulges have $R^{1/4}$ surface brightness profiles \\citep{dev48} while disks have exponential surface brightness profiles. As one goes from early-type to late-type galaxies one goes from galaxies that are bulge-dominated to galaxies that are disk-dominated. While this simplistic interpretation of the Hubble sequence has definite value, reality is considerably more complicated.  In recent years, our views of the Hubble sequence have evolved and gained more nuance. For elliptical galaxies it has become clear that they are not necessarily pure bulge systems: many elliptical galaxies contain embedded disks.  There is evidence from other information (e.g., kinematics) that elliptical galaxies form a heterogeneous class of galaxies that may have formed in different ways \\citep[e.g.][]{kor96}. For spiral galaxies a clearer understanding has developed of their bulge properties. High-resolution imaging - both from the ground \\citep[e.g.][]{dej96a}, and with the {\\it Hubble Space Telescope} \\citep[HST,][]{car98} - has shown that the central surface brightness profile (SBP) of many late-type spirals cannot be fit by the classical $R^{1/4}$ law that is well suited to describe the bulge profiles of early-type spirals. Instead, the SBPs of many late-type spirals rise above the extrapolation of the exponential disk in a way that can be well described by a second exponential \\citep{and94}. This has led to the now popular view that spiral bulges come in two flavors: on the one hand, the classical $R^{1/4}$ bulges which are mostly observed in early-type spirals, and on the other the ``pseudo-bulges'' \\citep{kor93} or ``exponential bulges'' \\citep[][]{car98} which are prevalent in later Hubble types. In reality there is probably a continuum of properties, instead of a dichotomy. When $R^{1/n}$ profiles \\citep{ser68} are fit to available SBPs, the profile shape parameter spans the full range of values $1 \\leq n \\leq 6$; the profile shape parameter correlates with both Hubble type and bulge-to-disk ratio of the galaxy\\footnote{Obviously, Hubble type and bulge-to-disk ratio are themselves strongly correlated through the definition of the Hubble sequence.}, in the sense that spiral galaxies with earlier Hubble type have bulges with higher $n$ values \\citep*{and95,gra01}.  The existence of different types of bulges in disk galaxies can be plausibly explained in the context of popular scenarios for the formation and secular evolution of galaxies. The classical massive $R^{1/4}$ law bulges fit in with the ``primordial collapse'' formation scenario first suggested by \\cite*{egg62}, in which the bulge forms during the initial collapse of a galaxy-sized density perturbation, and later ``acquires'' a disk through accretion processes. By contrast, the pseudo-bulges may have formed by secular evolution of a pre-existing disk, so that they formed after the disk, out of disk material. Some support for this scenario comes from the fact that pseudo-bulges are dynamically similar to their host disks \\citep{kor93}. Plausible secular evolution scenarios include the accretion of satellite galaxies \\citep{agu01}, buckling instabilities in a stellar bar \\citep*{rah91}, and the disruption of a stellar bar through the accumulation of a central mass concentration \\citep*{nor96}. Many discussions of these and related topics can be found in the review by \\cite*{wys97} and in the proceedings of the recent workshop on `The Formation of Galactic Bulges' \\citep*{car99}.  In the present paper we study the presence and properties of bulges in the very latest-type spiral galaxies (Scd or later). This is an important topic for several reasons. First, these galaxies are generally classified as very late type spirals because they do not have a very prominent bulge. As a result, many observational studies of bulges have avoided these galaxies. Second, it has become clear from recent work with HST that the majority of spiral galaxies contain a central star cluster. In the very latest-type spiral galaxies we find that $\\sim 75$\\% of the galaxies contain such a star cluster \\citep[][hereafter paper I]{boe02}. In late Hubble types, these clusters are easily mistaken for a small bulge when observed from the ground, even in good seeing conditions. So the bright, compact ``bulges'' in late-type spirals which were used as a classification criterion in the original work of \\cite{hub26,hub36} may in fact be dense star clusters occupying the photocenter of the galaxy. The purpose of this paper is to shed some light on these issues. In particular, we investigate whether the very latest-type spirals are completely bulgeless, whether they show excess light above the constant scale-length disk, and if so, whether this in fact implies the presence of a separate entity which could rightfully be called a bulge.  HST resolution is needed to separate the luminous nuclear star cluster from a putative bulge. Our I-band snapshot survey of late-type spiral galaxies conducted with the {\\it Wide Field and Planetary Camera 2} (WFPC2) and discussed in Paper~I therefore forms the basis of our analysis. We complement the HST observations with ground-based data that extends to larger radii. The paper is organized as follows: in \\S~\\ref{sec:data}, we describe the data and the analysis methods that form the basis of our work.  The results of our analysis are summarized in \\S~\\ref{sec:results}.  We discuss the implications of our findings, and present our conclusions in \\S~\\ref{sec:disc}. ",
        "conclusions": "}  We have presented an investigation into the structural properties of 19 spiral galaxies with Hubble type between Scd and Sm. From a combination of our high-resolution HST data and wide-field ground-based images, we obtain surface brightness profiles (SBPs) for the sample galaxies that cover a large dynamic range in galactic radius. We use these profiles to study quantitatively the presence and properties of bulges in spiral galaxies of the very latest Hubble types. Of course, by the very definition of the Hubble sequence, one does not expect very prominent bulges in these galaxies. However, it has been somewhat of an open question whether bulges are present at all in these galaxies. It has been realized only recently that space-based resolution is required to properly address this question. The nuclear morphologies of most spiral galaxies are complex, and a large fraction of spiral galaxies has a nuclear star cluster. These clusters are easily mistaken for ``compact bulges'' when observed with ground-based resolution. Previous studies with the HST have focused in majority on earlier Hubble types, and the present study is the first to focus exclusively on the very latest Hubble types.  Approximately 30\\% of the sample galaxies seem to be more-or-less ``pure'' exponential disks, without any type of a central bulge. Our sample was selected to focus on face-on galaxies, but studies of edge-on galaxies support the view that disk galaxies can indeed be completely bulgeless. \\cite*{mat99} and \\cite*{mat00} have studied edge-on ``super-thin'' galaxies such as UGC\\,7321, and have demonstrated that in these galaxies there is no evidence for a spheroidal component.  Despite being disk-dominated systems, most galaxies in our sample have SBPs that cannot be well fit by a single exponential, in the sense that the surface brightness in the central few kpc exceeds the inward extrapolation of the outer exponential disk. This has generally been found for other samples of spirals as well, and has generally prompted the addition of a bulge component to analytic models of the SBP. In particular, numerous studies have shown that the SBPs of intermediate- to late-type spirals can be well fit by a sum of two exponentials. The inner exponential has traditionally been interpreted as an `exponential' bulge or `pseudo' bulge. Such bulges can be (qualitatively) explained theoretically as a result of secular evolution of the disk. While this may be correct, we point out that this is not an unambiguous interpretation. In the absence of information on three-dimensional structure or dynamics there is no guarantee that one is dealing with a bona-fide bulge component. We explicitly illustrate this point for the galaxies in our sample that are not well fit by a single exponential. Indeed, the SBPs of these galaxies are generally well fit by a sum of two exponentials. However, most profiles can be described at least equally well with a S\\'ersic-type $R^{1/n}$ model over the entire radial range (outside the nuclear star cluster). The shape parameter $n$ is in the range $1 \\lta n \\lta 2.5$, which is not unrealistically large. As we have discussed in \\S\\ref{subsec:really}, there is no a priori theoretical reason for the SBPs of disk galaxies to be pure exponentials; S\\'ersic-type profiles are in some sense equally arbitrary as the model of choice. So it may well be that we are simply dealing with non-exponential disks in most of our sample galaxies, and that bona-fide bulges are rare at these Hubble types.  We have found that a number of late-type galaxies show central excess emission on spatial scales of a few hundred parsec that can be attributed neither to the nuclear cluster, nor to the S\\'ersic-type description of the stellar disk, nor to what one would generally consider to be a bulge component. The origin of this light component remains unclear.  One of the interesting findings from our work is that, despite the narrow range in Hubble type, the SBPs of the sample galaxies are far from uniform. Our study finds no systematic trends in the structural properties with morphological type. The exact Hubble type of spirals between Scd and Sm appears somewhat arbitrary, in the sense that it provides little information about the presence and relative importance of galaxy bulges. It is quite possible that this may be due to the presence of nuclear clusters, which may have played an important role in the morphological classification in photographic catalogs. In general our results fit in well with the picture that emerged from a ground-based imaging study of 49 late-type spirals by \\cite{mat97}. They found that late-type spirals exhibit a diverse array of structural properties and morphologies, even in galaxies with otherwise similar parameters, and they concluded that bulges are often very weak or non-existent. In studies of this kind one does need to be concerned about selection effects and observational bias. The \\cite{mat97} sample was selected for low galaxy luminosity ($M_V \\gta -18.8$), and one might worry that it did not provide a representative view of the family of very late-type spirals. Our galaxy sample, on the other hand, has been selected only for Hubble type, distance, and inclination. It thus includes more luminous (and presumably more massive) galaxies than the \\cite{mat97} sample. The blue absolute luminosity of the galaxies in our sample ranges from $M_B = -17.4$ (NGC\\,2552, NGC\\,4701) to $M_B = -20.5$ (NGC\\,2805). However, we find no systematic correlation between absolute luminosity and position along the sequence of SBP shapes discussed in \\S\\ref{subsec:seq}. This suggests that more luminous late-type spirals are not systematically different from their faint-end relatives, and that presumably the \\cite{mat97} results are valid for most late-type spirals.  Late type spiral galaxies are in most ways ``normal'' spiral galaxies. Their angular momenta and rotation velocities are not atypical, as demonstrated by a recent survey of optical rotation curves \\citep{mat02}. They are in some sense the dynamically simplest type of disk galaxies. They are (mostly) disk-dominated and often have only faint spiral arm structure or even no detectable density perturbations at all. Yet, we are clearly a long way from understanding their formation and evolution. The latest-type disk galaxies provide the most stringent observational constraints on the well-known angular momentum problem in cold dark matter (CDM) galaxy formation models \\citep[e.g.][]{nav00}. Because of their unevolved disks and apparent history of only modest star formation, they are a challenge for proposed solutions for the angular momentum problem which rely on energy feedback from supernova explosions \\citep{sil01}. To gain a better understanding of these issues it will be important to continue to improve our knowledge of the structure of late-type spiral galaxies. Observational studies such as the one presented here will continue to be essential."
    },
    "0212/astro-ph0212308_arXiv.txt": {
        "abstract": "We present an analysis of {\\it BeppoSAX}  observations of the IC1262 galaxy cluster and report the first temperature and abundance measurements, along with preliminary indications of diffuse, nonthermal emission.  By fitting a 6$\\arcmin$ ($\\sim$360 {\\it h$_{50}^{-1}$} kpc) region with a single Mewe-Kaastra-Liedahl model with photoelectric absorption, we find a temperature of 2.1 - 2.3 keV, and abundance of 0.45 - 0.77 (both 90\\% confidence).  We find the addition of a power-law component provides a statistically significant improvement ({\\it F}-test = 90\\%) to the fit.  The addition of a second thermal component also improves the fit but we argue that it is physically implausible.  The power-law component has a photon index ($\\Gamma_{X}$) of 0.4 - 2.8 and a nonthermal flux of (4.1  - 56.7) $\\times$ 10$^{-5}$ photons cm$^{-2}$ s$^{-1}$ over the 1.5 - 10.5 keV range in the Medium Energy Concentrator spectrometer  detector.  An unidentified X-ray source found in the {\\it ROSAT} High Resolution Imager observation ($\\sim$0$\\arcmin$.9 from the center of the cluster) is a possible explanation for the nonthermal flux; however, additional evidence of diffuse, nonthermal emission comes from the NRAO VLA Sky Survey and the Westerbork Northern Sky Survey radio measurements, in which excess diffuse, radio flux is observed after point-source subtraction.  The radio excess can be fitted to a simple power law with a spectral index of $\\sim$1.8, which is consistent with the nonthermal X-ray emission spectral index.  The steep spectrum is typical of diffuse emission and the size of the radio source implies that it is larger than the cD galaxy and not due to a discreet source. ",
        "introduction": "Since the discovery of nonthermal emission in galaxy clusters, in the form of a radio halo in Coma \\citep{will}, there has been a debate about the nature of its origin.  \\citet{jaf} demonstrated that although Active Galactic Nuclei (AGNs) in clusters could explain the numbers of relativistic electrons observed, this model could not explain their ability to diffuse through the entire cluster within their short ($\\sim$10$^{8}$ yr) radiative lifetime.  \\citet{hol} proposed that primary electrons are able to diffuse at velocities greater than the Alfv\\'{e}n velocity and therefore could diffuse from point sources throughout the ICM within their radiative lifetime.  There has also been some discussion that secondary electrons could explain the diffuse relativistic electron  population.  Secondary electrons are produced in proton-proton interactions from a population of relativistic protons \\citep{den}, which have a relatively long radiative lifetime ($\\sim$10$^{10}$ yr) \\citep{blasi01}.  The majority of authors, however, have looked to in-situ re-acceleration models from MHD turbulence or shocks \\citep[for example]{eil} to explain diffuse nonthermal emission from clusters.  Mergers are the most popular theory for the origin of these shocks, since they provide the energy needed to produce a radio halo \\citep{trib,burns01,burns02}, and are able to transport electrons via bulk flows \\citep{roe}.  \\citet{reph} correlated synchrotron emission to X-ray emission by demonstrating that cosmic microwave background (CMB) photons will inverse Compton scatter (ICS) off relativistic electrons, and emerge in the X-ray regime.  The relationship between the X-ray and radio flux is determined by the strength of the cluster magnetic field.  With only $\\sim$30 clusters known to have radio halos \\citep[and references therein]{giov02}, and three hot, rich clusters: Coma \\citep{fusco01}, A2256 \\citep{fusco02}, and A2199 \\citep{kaa01}, and one group: HCG 62 \\citep{fuka}, with detected nonthermal X-ray emission, it is difficult to constrain the current models.  Detection of nonthermal emission in hot clusters is complicated because the X-ray emission is dominated by thermal emission in the 0.5 to 25 keV energy band, which is why the Phoswich Detection System (PDS) has been required to make nonthermal X-ray detections in hot clusters.  In cool clusters, such as IC1262 ({ \\it kT} $\\sim$2keV), any nonthermal emission should be detectable in the 1.5 - 10.5 keV range and observable in the Middle Energy Concentrator Spectrometer (MECS), which unlike the PDS, has spatial capabilities.  Therefore, cool clusters provide an opportunity to add additional constraints, because the spatial information from the MECS can be used to constrain the regions of nonthermal emission, and reduce the chances of contamination by AGNs.  The IC1262 galaxy cluster (also Zwicky Cluster 1728.5+4353) is a  poor cluster of galaxies at a red-shift of z = 0.0343 \\citep{col}, so that an angular distance of 1$\\arcmin$ corresponds to 60 {\\it h$_{50}^{-1}$} kpc.  The cluster is so-named because it is dominated by the X-ray bright cD galaxy, IC1262 (z = .0328 \\citep{weg}).  The cluster was observed by the HRI in March 1997 for $\\sim$26 ksec, and {\\it BeppoSAX} in February 1999 for $\\sim$100 ksec.  \\citet{trin} analyzed the HRI data, mainly focusing on the cD galaxy.  Throughout this letter, we assume a Hubble Constant of H$_{0}$ = 50 {\\it h$_{50}$} km s$^{-1}$ Mpc$^{-1}$ and q$_{0}$ = $\\case{1}{2}$.  Quoted confidence intervals are at a 90\\% level, unless otherwise specified. ",
        "conclusions": "IC1262 is a poor, cool cluster that shows some evidence of nonthermal X-ray emission. While, a similar claim can be made for another low mass system, the compact group HCG 62, IC1262 is distinguished in that it shows evidence for a radio halo, where HCG 62 does not \\citep{fuka}. Currently, cluster mergers are the most popular model for the cosmic ray acceleration that leads to the production of diffuse nonthermal emission. IC1262 shows no direct X-ray evidence of a merger, such as elongated X-ray emission (as in the case of A754), or bi-modal X-ray morphology (as in the case of A1750).  There is, however, optical evidence that the cluster may have undergone a merger.  The cD galaxy has a red-shift of z = 0.0328 (9858 km s$^{-1}$)\\citep{weg} compared to a red-shift of z = 0.0343 (10311 km s$^{-1}$) \\citep{weg} for the cluster.  This gives a peculiar velocity, $\\Delta$v, of 453 km s$^{-1}$, which is $>$3$\\sigma$ above the dispersion of peculiar velocities (168$^{+41}_{-34}$ km s$^{-1}$) of cD galaxies found by  \\citet{oeg}.  A line of sight merger could be the explanation for this unusually high peculiar velocity.  In addition, \\citet{trin} analyzed the {\\it ROSAT} HRI observation of IC1262 and found bright arc near the cD galaxy, which may indicate dynamic evolution of the cD galaxy due to a merger.  Our analysis has suggested that diffuse, nonthermal emission is associated with IC1262.  We find a 2.78$\\sigma$ detection in the MECS, but no detection in the LECS within 90\\% confidence.  This lack of detection in the LECS is probably because it is not as sensitive to hard photons as the MECS.  Although our 90\\% confidence in $\\Gamma_{X}$ and normalization indicate that U1275\\_09534024 could be the reason for our nonthermal detection, the radio implies the existence of diffuse nonthermal emission.  When the diffuse radio excess is fit to a simple power law, it gives a spectral index, $\\alpha_{r}$ $\\sim$1.8.  When we use the radio to confine $\\Gamma_{X}$ to 2.8, then the nonthermal, X-ray emission can not be accounted for with point sources.  We also reiterate that U1275\\_09534024 may not even be emitting nonthermal, X-ray emission. A deeper radio observation will be able to confirm the existence of a radio halo and provide a spectral index that can be used to constrain $\\Gamma_{X}$.  With a sensitivity up to $\\sim$30 keV, the {\\it Rossi X-Ray Timing Explorer} Proportional Counter Array(PCA) should be able to detect any diffuse, nonthermal, X-ray emission associated with IC1262, and also would be able to determine the variability, if any, of U1275\\_09534024."
    },
    "0212/astro-ph0212144_arXiv.txt": {
        "abstract": "OGLE and ASAS are long term observing projects operated by the Warsaw University Observatory at the Las Campanas site in Chile.  OGLE is currently monitoring almost 200 million stars in the Galactic Bulge and the Magellanic Clouds, and has detected so far almost 1,000 events of gravitational microlensing with the dedicated 1.3-meter telescope.  ASAS uses several very small instruments to monitor all southern sky for variability down to approximately 14 magnitude.  A total of almost 300 thousand variable stars were discovered so far by the two projects, and all photometric data is available on the WWW.  Both projects aim at real time recognition and verification of all new phenomena in the sky.  OGLE is likely to discover planets by 2003 and stellar mass black holes by 2004-2005.  All OGLE and ASAS data is made public domain as soon as possible, and may be used by a Virtual Observatory. ",
        "introduction": "OGLE (Optical Gravitational Lensing Experiment, the leader: Andrzej Udalski) is a long term project, with a gradual expansion of its capability. The dedicated 1.3 meter telescope is located at the Las Campanas Observatory in Chile, a site owned and operated by the Carnegie Institution of Washington.  OGLE is operated by the Warsaw University Observatory.  Information about it may be found at: \\vskip 0.2cm \\centerline{OGLE: \\hskip 1.0cm http://sirius.astrouw.edu.pl/\\~~ogle/} \\centerline{OGLE: \\hskip 1.0cm http://bulge.princeton.edu/\\~~ogle/ \\hskip 0.05cm} \\vskip 0.2cm  OGLE-III is the current, third stage of the project.  The observations are done with a mosaic $ 8K \\times 8K $ CCD camera built by A. Udalski.  OGLE-III begun in May 2001.  In the 2002 Galactic Bulge season almost 400 candidate microlensing events were discovered:  \\centerline{OGLE-EWS: \\hskip 1.0cm http://sirius.astrouw.edu.pl/\\~~ogle/ogle3/ews/ews.html} \\noindent A search for planetary transits was conducted on 32 nights during a 45 day interval in 2001.  A total of 59 stars were identified, for which low depth flat bottom transits with orbital periods shorter than 10 days were found (Udalski et al. 2002a,b).  Some of these are likely caused by `hot Jupiters', others by brown and red dwarfs.  The analysis of OGLE-II stage (Udalski, Kubiak \\& Szyma\\'nski 1997), which covered years 1997-2000, is advanced, and almost all data is in public domain.  This includes catalogs of I, V, B magnitudes and positions for a total of almost 40 million stars in the Galactic Bulge, and in the Magellanic Clouds (Udalski et al. 2002c, 2000b, 1998a), catalogs of over 500 microlensing events (Udalski et al. 2000a, Wo\\'zniak et al. 2001), catalogs of over 270 thousand variable stars (Zebru\\'n et al. 2001, Wo\\'zniak et al. 2002), and in particular thousands of eclipsing binaries in the SMC (Udalski et al. 1998a) and Cepheids in the LMC and SMC (Udalski et al. 1999a,b).  In addition OGLE astrometry provided proper motion measurements of thousands of stars (Soszy\\'nski et al. 2002, Sumi et al. 2002), leading to the discovery of streaming motion (rotation) in the Galactic Bar.  Of particular interest are special microlensing events.  OGLE-2000-BUL-43 was found to have a spectacular parallax effect (Soszy\\'nski et al. 2001). Even more dramatic was OGLE-1999-BUL-19, the first event with multiple peaks in its apparent brightness caused by the Earth's orbital motion and a very long time scale: $ t_E = R_E / V = 372 $ days (Smith et al. 2002). Even longer time scales were found for OGLE-1999-BUL-32 (Mao et al. 2002, $ t_E = 641 $ days and OGLE SC\\_5 2859 (Smith 2002, $ t_E = 551 $ days). The last two events are likely caused by massive lenses, probably stellar mass black holes. ",
        "conclusions": ""
    },
    "0212/astro-ph0212528_arXiv.txt": {
        "abstract": "{Nayakshin \\& Kazanas (2002) have considered the time-dependent illumination of an accretion disc in Active Galactic Nuclei, in the lamppost model, where it is assumed that an X-ray source illuminates the whole inner-disc region in a relatively steady way. We extend their study to the flare model, which postulates the release of a large X-ray flux above a small region of the accretion disc. A fundamental difference to the lamppost model is that the region of the disc below the flare is not illuminated before the onset of the flare. After the onset, the temperature and the ionization state of the irradiated skin respond immediately to the increase of the continuum, but pressure equilibrium is achieved later. A few typical test models show that the reflected spectrum that follows immediately the increase in continuum flux should always display the characteristics of a highly illuminated but dense gas, i.e. very intense X-ray emission lines and ionization edges in the soft X-ray range. The behaviour of the iron line is however different in the case of a ``moderate\" and a ``strong'' flare: for a moderate flare, the spectrum displays a  neutral component of the Fe K$\\alpha$ line at 6.4 keV, gradually leading to more highly ionized lines. For a strong flare, the lines are already emitted by FeXXV (around 6.7 keV) after the onset, and are very intense, with an equivalent width of several hundreds eV. A strong flare is also characterized by a steep soft X-ray spectrum. The variation timescale in the flare model is likely smaller than in the lamppost model, due to the smaller dimension of the emission region, so the timescale for pressure equilibrium is long compared to the duration of a flare. It is therefore highly probable that several flares contribute at the same time to the luminosity. We find that the observed correlations between $R$, $\\Gamma$, and the X-ray flux are well accounted for by a combination of flares having not achieved pressure equilibrium, also strongly suggesting that the observed spectrum is always dominated by regions in non-pressure equilibrium, typical of the onset of the flares. Finally, a flare being confined to a small region of the disc, the spectral lines should be narrow (except for a weak Compton broadening) and Doppler shifted, as stressed by Nayakshin \\& Kazanas (2001).  All these features should constitute specific variable signatures of the flare model, distinguishing it from the lamppost model. It is however difficult, on the basis of the present observations and models, to conclude in favor of one of the hypothese. ",
        "introduction": "The model of ``irradiated accretion disc'' is now widely accepted as an explanation for the presence in the X-ray spectrum of Active Galactic Nuclei (AGN) of a Fe K$\\alpha$ line and an X-ray hump (Pounds et al. 1990), and for the absence of time delay (or for the very short delay) between the variations of the optical and UV continuum fluxes (Collin-Souffrin 1991, Clavel et al. 1992). Both phenomena are most likely produced by an X-ray flux irradiating a cold medium (the accretion disc?), which leads to a signature as reprocessed radiation in the UV range, and to a Fe K$\\alpha$ line and a Compton reflection hump in the X-ray range.  Three different models have been proposed to account for these properties: an X-ray source (whose origin is unknown) at a relatively high altitude above the disc, illuminating a large central region (the ``lamppost''), a hot corona completely covering a cold disc (Haardt \\& Maraschi 1991, 1993), and finally a ``patchy corona'' (Haardt et al. 1994) partially covering the disc.  This last model is often preferred for various reasons: the shape of the reflected continuum, the correlation between the Fe K$\\alpha$ and the continuum, the absence or weakness of the helium- and hydrogen-like Fe lines (cf.  Nayakshin 2000, Nayakshin \\& Kallman 2001), the fact that it allows us to account for the large UV over X observed luminosity ratio (cf.  Walter et al. 1994). It is called the ``magnetic flare'' model, as it refers to magnetic loops reconnecting and suddenly releasing their energy as X-ray photons on a small portion of the disc.  There is presently no clear consensus concerning the physics of the flare, which is generally considered to be analogous to the sun. In particular, the height of a flare above the disc, its luminosity and its compactness are free parameters in addition to others that can be chosen in order to fit the observed variational and spectral properties of the objects. Poutanen \\& Fabian (1999) and Merloni \\& Fabian (2001) have proposed a model accounting for the spectral variability of accreting black holes (it is scaled for galactic sources, but it applies to AGN as well). Many simultaneous flares are required to account for the observed luminosity of an AGN. To avoid the smearing of strong luminosity variations, they propose that the amplitude of variability can be explained if these flares are triggered in a chain reaction. Each flare produces an ``avalanche'' of neighbouring flares, whose duration and amplitude determine the timing properties. In particular, the distribution of size and luminosity of the avalanches accounts for the shape of the power spectrum density, which gives the power as a function of the Fourier frequency  (PSD, see also Zycki 2002). The most powerful avalanches have the largest size (corresponding to the largest time scales), and they correspond to the ``break'' and the red branch of the PSD (cf.  below).  In this model, the height of a magnetic loop increases during the avalanche, increasing the number of UV photons available for the cooling of the corona and causing the observed softening of the X-ray spectrum. It should be at least one order of magnitude larger than the pressure scale height of the disc. Consequently an avalanche can illuminate a relatively large fraction of the inner disc (since the radius of the illuminated region is of the order of the height of the source). Nayakshin \\& Kazanas (2001) argue on the contrary that the height of a flare should be of the order of the scale height of the disc, and that an avalanche of flares should illuminate only a very small fraction of the disc. In both cases the total number of simultaneous powerful avalanches is small, so that the total area illuminated by them is small compared to the disc area.  This is a fundamental issue as {\\it it implies that the region of the disc located under an avalanche was not illuminated before the onset of the avalanche}.  Our aim in this paper is to study the implication of this model, in particular on the soft X-ray spectrum and on the Fe lines. Note that in the following we will use the word ``flare\", even if it is actually an avalanche.  Several recent results concerning time variations of the Fe lines are very intriguing in the context of irradiated discs. The expectation would be that the Fe lines respond to variations of the X-ray continuum flux in a simple way, in the same way as the broad optical-UV lines respond to the UV-X continuum flux. This is far from being the case, and the behaviour of the lines is actually very complex. The interpretations of the results are sometimes even contradictory, for instance for MCG -6-30-15 (Lee et al. 1999, Vaughan \\& Edelson 2001), where it is not clear whether the line flux stays constant, while the continuum varies.  Another example is Mkn 841, where a rapid  (less than 15 hours) and strong (a factor two) variation of the Fe K$\\alpha$ line is related only to a weak variation (less than 20 $\\%$) of the X-ray continuum (Petrucci et al. 2002).  Moreover the variations of the optical/UV continuum, which are generally assumed to be induced by variations of the X-ray flux and should therefore follow them after a short time delay if they are emitted by the accretion disc (Rokaki et al. 1993, Kazanas \\& Nayakshin 2001), do not obey a simple relationship with the X-ray light curves. For instance in NGC 7469, the UV flux leads the 2-10 keV flux in the peaks of luminosity (Nandra et al. 1998) while the 2-10 keV flux leads the UV flux during the minima of luminosity. However this effect can be an artefact due to the limited range of energy considered: since the X-ray spectrum becomes softer when the UV flux is large, the X-ray flux might be reprocessed below 2 keV. In NGC 3516 there is no clear correlation between the UV and the X-ray flux (Edelson et al. 2000).  It is clear that the response of the irradiated medium is not simply the consequence of the variation of the ionization state  and of the thickness of the ``hot skin'' of the disc irradiated by a variable X-ray flux, as discussed either for a constant density medium or for a medium in hydrostatic equilibrium (Zycki \\& R\\'o\\.za\\'nska 2001). A most interesting idea stressed for the first time by Nayakshin \\& Kazanas (2002, hereafter NK02) is that if the illuminating flux varies on short time scales, the Fe line flux should not be a function of the instantaneous illuminating spectrum. In response to a variation of the illuminating flux, the atmosphere of the accretion disc remains  during a relatively long time, required for readjustment of the hydrostatic balance, in a non-equilibrium state.  In this paper we extend this idea, discussing the fate of a region of the disc illuminated by a flare, a case not considered by NK02, who limited their study to the ``lamppost model''. To clarify the discussion, Fig. \\ref {fig-lamppost-flare} shows our vision of the lamppost and the flare models.  \\begin{figure} \\begin{center} \\psfig{figure=fig-lamppost.eps,width=9cm} \\psfig{figure=fig-flare.eps,width=9cm} \\caption{ The lamppost and the flare models. } \\label{fig-lamppost-flare} \\end{center} \\end{figure}  In both cases, the disc is illuminated by an X-ray source whose luminosity is equal to a fraction of the bolometric luminosity, located at a given height $H$ above the disc. The lamppost is located in the center (above the black hole), and its height is of the order of 10-20 $R_{G}$ ($R_{G}$ being the gravitational radius $GM/c^2$), while a flare is located at any radius, and its height above the disc is much smaller. Several flares can be shining at the same time. The dimension of the region illuminated by a source located above the disc is of the order of its height, and the flux on this region is of the order of $L/4\\pi H^{2}$. For a flare, the flux is thus much larger than for a lamppost. It is also much larger than the viscous flux provided by the release of gravitational energy, which is maximum at a few $R_{G}$ and decreases as $R^{-3}$ (where $R$ is the distance to the center).  According to this discussion, we assume therefore that the flare model differs from the lamppost model in three respects:  \\begin{itemize}  \\item In the lamppost model, the disc is illuminated permanently. The illuminating flux can switch from a state $F_{X1}$ to a state $F_{X2}$, with $F_{X2}\\sim$ a few $F_{X1}$, or inversely, but the illumination is {\\it always} present.  Therefore an irradiated layer, called the ``hot skin'', always exists above the inner disc. However the hydrostatic equilibrium corresponding to the two states is different, and the hot skin as well as its spectrum will evolve slowly (cf. later) from one equilibrium state to the other.  On the contrary, in the flare model, there is {\\it no illumination} in the initial state (and in the final state, after the extinction of the flare), so the newly formed hot skin is much denser than an illuminated atmosphere (it does not exclude that another part of the disc is already illuminated, as we will discuss in Section 4).  \\medskip \\item In the lamppost model, the disc is illuminated by a flux $F_X$ smaller than or of the order of the viscous flux $F_d$, while in the flare model, $F_X$ is much larger than $F_d$.  \\medskip \\item In the lamppost model, the whole inner disc is illuminated, while in the flare model the illumination acts only on a small fraction of the disc.  \\end{itemize}   In the following section, we recall the relevant time scales, and in Section 3 we give results corresponding to a few typical cases. Implications for the observations are discussed in Section 4. ",
        "conclusions": "We have made a preliminary study of the consequences of the flare model on the X-ray spectrum and on its time variations.  We have in particular stressed the importance of the transient state during which the flare reached its maximum intensity and the irradiated atmosphere below has adjusted its thermal and ionization equilibrium, but is not in pressure equilibrium.  Such states are characterized by intense soft X-ray lines, and in the case of strong flares, by a very intense FeXXV line at 6.7 keV with a neutral component at 6.4 keV.  We have also reached the conclusion that the observed correlation between $F$(2-10keV) and $\\Gamma$ can be explained without invoking a feed-back of the reprocessed variation on the primary one. However an individual flare alone is not able to provide this correlation, as the simple variation of albedo occuring during the transition towards pressure equilibrium leads to a correlation inverse to the observed one, and in particular to a steep soft X-ray spectrum. The onset of a strong flare superposed on several preexisting flares is required. This strongly suggests that the luminosity is dominated by transient states having not reached pressure equilibrium. On the other hand the correlation between $\\Gamma$ and $R$, if real, is easily accounted for by all models.  In the absence of pressure equilibrium, such atmospheres are not subject to the usual thermal instability, that leads to the disappearance of the intermediate temperature and intermediate states of ionization. As a consequence the spectrum displays both intense neutral and ionized iron lines. Quite strangely, it would mean that the first models proposed for irradiated discs, and assuming a constant density illuminated slabs (Ross \\& Fabian 1993, and susbsequent papers) are perhaps more relevant than the more sophisticated ones with atmospheres in hydrostatic equilibrium!  In this study, we have not taken into account any relativistic effects, and we have considered only {\\it local} profiles. This is justified by the fact that the emission is provided by small localized regions on the disc. Thus the line profiles are not broadened by relativistic and gravitational effects, but only by Comptonization (which actually might be important, see the figures of the paper). On the other hand, the lines may be boosted, and red or blueshifted, and they could move according to the orbital motion of the disc, if the emission region is located close to the black hole.  \\medskip  Several aspects should therefore differentiate the flare from the lamppost model: \\begin{itemize}  \\item a shorter variation time scale, of $10^{3}-10^{4}$ s,   \\item larger spectral variations during the increase of the  flux, including the appearance of lines and edges in the soft X-ray range and of very intense Fe K$\\alpha$ lines at 6.4 and 6.6 keV,  \\item narrow and moving lines.  \\end{itemize}    We are not able to reach definitive conclusions about the reliability of either the lamppost or the flare model in the context of present variation studies, though a few observations tend to support the flare model. More observations are required to understand the relationship between the flux, the iron line, the soft and hard X-ray continua. To make some progress, it is also necessary to compare a larger set of models to the observations.  The aim of this paper was simply to draw  attention to the differences between the lamppost and the flare model, in particular to the fact that the disc underlying a flare is ``cold\" before the flare settles, and more generally to the fact that the luminosity is certainly dominated by transient states out of pressure equilibrium. We have given here only a few  typical examples, postponing a more in-depth study to a future paper."
    },
    "0212/astro-ph0212002_arXiv.txt": {
        "abstract": "\\noindent We construct a physically motivated analytical model for the quasar luminosity function, based on the joint star formation and feeding of massive black holes suggested by the observed correlation between the black hole mass and the stellar mass of the hosting spheroids. The parallel growth of massive black holes and host galaxies is assumed to be triggered by major mergers of haloes. The halo major merger rate is computed in the frame of the extended Press-Schechter model. The evolution of black holes on cosmological timescales is achieved by the integration of the governing set of differential equations, established from a few reasonable assumptions that account for the distinct (Eddington-limited or supply-limited) accretion regimes. Finally, the typical lightcurves of the reactivated quasars are obtained under the assumption that, in such accretion episodes, the fall of matter onto the black hole is achieved in a self-regulated stationary way. The predicted quasar luminosity function is compared to the luminosity functions of the 2dF QSO sample and other, higher redshift data. We find good agreement in all cases, except for $z<1$ where the basic assumption of our model is likely to break down. ",
        "introduction": "\\label{intro}  In the last decade, a considerable progress has been achieved towards the understanding of the nature and evolution of Active Galactic Nuclei (AGN). It is now commonly accepted that they are powered by Massive Black Holes (MBHs) fuelled by material originating in the hosting dark-matter haloes and infalling from the close galactic environment in the process described as accretion. New models connecting the AGN phenomenon to the hierarchical growth of haloes have been developed, based on increasingly robust theoretical background substituting phenomenology. In this frame, the masses of MBHs are usually assumed to scale with the masses of the hosting dark-matter haloes with density estimated from the Press \\& Schechter (1974) formalism (\\eg \\citealt{efstathiou88,haiman98}). The AGN luminosity is supposed close to the Eddington value during a ``duty cycle'' that can be as short as $10^6$ to some $10^7$ years (\\eg \\citealt{haiman98,martini01, wyithe02}) or longer, of the order of some $10^8$ years (\\eg \\citealt{haehnelt93,haehnelt98}). With time, these models have become more and more complex and complete. The accretion rate onto the MBH varies depending on the model: it can be constant or redshift-dependent (\\citealt{kauffmann00,haiman00}). The AGN activity in many cases appears to be a recurrent event (\\citealt{siem97,hse01}). The diverging opinions as to this recurrence have converged, accounting for tidal interactions and merging effects. In addition, the Eddington-limited capability of MBHs to accrete mass and the progressive consumption of the hot gas available in haloes are also expected to play an important role in the evolution of AGN luminosities \\citep{cavitt00}.  In several recent papers (\\citealt{silkrees98,monaco00,granato01, archibald02,volonteri03} and references therein), the joint modelling of star formation and the growth of MBHs lying in galactic nuclei has been considered in order to account for the observed correlation between the masses of the MBH and the hosting spheroid \\citep{wandel99,woo02}. Spheroids grow through galaxy mergers, giving rise to new ellipticals, and through the mass transfer from unstable discs (possibly not even settled down) to bulges. Although none of these processes is directly associated with the Major Merger (MM) of dark-matter haloes, some concatenation of events is believed to correlate both phenomena. Firstly, MMs mark the beginning of the orbital decay, through dynamical friction, of the more massive galaxies of the progenitors, which eventually merge at the center of the new halo. Secondly, the shock-heated gas is redistributed and begins to cool from the central higher density region at the strongest rate ever reached during the whole halo life. Large amounts of gas then fall into the central galaxy making its disc easily become unstable (particularly at high redshifts where discs are more compact) and transfer mass into the bulge. As the spheroid grows some mass may reach the nuclear region and feed the central MBH. We therefore expect the growth of spheroids and MBHs not only go in parallel, but also be punctuated by the rate at which dark-matter haloes form as a consequence of MMs \\citep{kauffmann00,haiman00}.  It is difficult to tell at which extent this scenario is able to reproduce the observed properties of AGN and their dependence with redshift. The models found in the literature only address partial aspects of it and often include poorly justified assumptions and rough estimates of the MBH host density. Nowadays, the development of a complete model of this kind has become feasible. In particular, accurate analytical expressions for the halo MM rate are available (\\eg \\citealt{percival00,rgs01}) that can be easily implemented. The aim of the present paper is to build such a complete, physically motivated model and check the validity of the previous scenario by comparing the predicted luminosity function of quasars (QLF) with the most recent publicly available observations.  Our model is based on the three following ingredients: 1) an accurate expression, in terms of halo mass and cosmic time, for the halo MM rate (Sec.\\ \\ref{MMrate}) derived in the frame of the extended Press-Schechter model; 2) the evolving properties of MBHs (Sec.\\ \\ref{evoMBH}), obtained by solving analytically the governing set of differential equations established under a few simple assumptions consistent with the parallel growth of MBHs and their hosting galaxies from the hot gas available in the halo; and 3) the typical lightcurves of AGN associated with MBHs of different masses at different cosmic times, inferred under the assumption of an Eddington-regulated accretion rate (Sec.\\ \\ref{QLC}). By comparing the resulting theoretical QLFs (Sec.\\ \\ref{QLF}) with the ones observed at different redshifts, the best values of the free parameters entering the model are inferred. This leads to a fully consistent analytical expression of the QLF (Sec.\\ \\ref{data}) that recovers both the observed QLF shape and its redshift evolution at intermediate and high redshifts. In the last section, we discuss the possible shortcomings of the model and the implications of our results on the process of galaxy formation and evolution. Throughout the present work, we assume a $\\Lambda$CDM cosmology, defined by the values of the Hubble constant (in km s$^{-1}$ Mpc$^{-1}$), density parameter, cosmological constant, baryonic fraction, and normalisation of the power spectrum respectively given by $(H_0,\\Omega\\m,\\Omega_\\Lambda,\\Omega\\b h^2,\\sigma _8) = (70,0.3,0.7,0.02,0.75)$, with $h=H_0/100$. ",
        "conclusions": "\\label{discuss}  We have developed a fully consistent analytical model of QLF from reasonable physical assumptions and used available data on the QLFs at various redshifts to fix the values of the free parameters. The predicted QLF agrees with the available observations at redshifts higher than $\\sim 1$.  The fact that, the more accurate halo MM rate used in our model behaves essentially as the product of two PS mass functions instead of just a single one as often adopted explains the more marked positive evolution shown by the inferred QLF. Yet, the evolution so predicted is not strong enough to match the observations. The QLF calculated from the adopted MM rate but assuming a MBH mass proportional to the halo mass and the associated quasar radiating close to the Eddington regime fails to reproduce the observed data. A successful approach requires also modelling the typical evolution on cosmological timescales of MBHs and the typical lightcurve of AGN. In the present study, such a theoretical lightcurve rests on the assumptions of a phenomenological bell-shaped accretion rate and a physically grounded self-regulated radiation, while the modelling of the MBH growth is based on reasonable assumptions emanating from the expected connection between halo MMs and the reactivation of AGN.  \\begin{table*} \\caption{Best-fitting parameters, column 1, corresponding to four data sets. In columns 2 and 3, the observed QLFs drawn from the 2QZ survey in the two redshift intervals quoted. In column 4, the high-$z$ QLFs drawn essentially from the SDSS survey \\citep{fan01} and, in column 5, the integrated quasar light estimated from the absorption spectra of high-$z$ QSOs \\citep{mcdo01}. $z\\refe$ is the redshift at which the accretion regime changes from Edington to supply-limited. $k$ is the fractional mass increase of the MBH in the Edington-limited regime. $\\epsilon\\gal$ is the mass ratio between the MBH and the host galaxy and $\\epsilon\\g$ the ratio between the mass accreted by the MBH in a MM event and the total gas mass available, both in the supply-limited accretion regime. $\\tau\\acc$ is the typical timescale of accretion onto the MBH (essentially the quasar duty cycle). $M\\min$ and $M\\max$ are the lower and upper masses of haloes likely to harbour quasars. Finally, $\\nu$ is the fraction of MMs giving rise to the reactivation of one observed quasar (in the case of the integrated quasar light there is no value for $\\nu$ since both the predictions and observations are normalised at $z=2$). Errors correspond to $1\\sigma$ intervals, calculated in the standard way for a multiparametric fit by $\\chi^2$ minimization. The lack of errors in the second column is due to the poor fit obtained in this case, which makes such an estimation meaningless.}  \\label{tabres} \\begin{center} \\begin{tabular}{lllll} \\hline \\hline &~~~~~~2QZ &~~~~~~~~2QZ &~~~high-$z$ &~~~~~$\\rho_L$\\\\ & $0.3<z<2.3$                  &   $1.27<z<2.3$        & $3<z<4.7$     & $0<z<6$\\\\ \\hline \\hline $z\\refe$                        & $~~~~~1.92$  & $~~2.01^{+0.02}_{-0.02}$  & $~4.6^{+\\infty}_{-4.6}$    & $~1.97^{+0.04}_{-0.05}$\\\\ $\\epsilon\\gal~(\\times 10^{-3})$ & $~~~~~7.91$ &$~~2.45^{+0.02}_{-0.05}$ & $~6.6^{+6.7}_{-2.6}$ &  $~2.3^{+\\infty}_{-2.3}$\\\\ $\\epsilon\\g~(\\times 10^{-2})$   & $~~~~~5.08$  & $~~1.54^{+0.18}_{-0.18}$ & $~11.0^{+\\infty}_{-10.5}$  & ~$8.5^{+\\infty}_{-8.5}$\\\\ $\\tau\\acc$  (Gyr)               & $~~~~~0.05$  & $~~0.18^{+0.01}_{-0.02}$  & $~0.001^{+0.005}_{-0.001}$ & $~0.02^{+0.06}_{-0.02}$\\\\ $k$                             & $~~~~~9.21$  & $~~2.36^{+0.10}_{-0.09}$  & $~1.0^{+4.1}_{-0.5}$    & $~2.01^{+2.20}_{-0.17}$\\\\ $\\log M\\min$ (M$_\\odot$)        & $~~~~~10.72$ & $~~12.47^{+0.01}_{-0.01}$  & $~11.4^{+0.2}_{-\\infty}$   &$~10.0^{+0.7}_{-1.2}$\\\\ $\\log M\\max $ (M$_\\odot$)       & $~~~~~12.88$ & $~~13.72^{+0.09}_{-0.04}$ & $~14.6^{+\\infty}_{-2.0}$  &$~13.50^{+\\infty}_{-0.6}$\\\\ $\\nu $                          & $~~~~~0.02$ & $~~0.11^{+0.01}_{-0.01}$ & $~0.05^{+0.07}_{-0.03}$  &  \\,\\,\\,~~~~--  \\\\ \\hline \\hline \\end{tabular} \\end{center} \\end{table*}  The different samples of quasars studied with redshifts ranging from 1.27 to 6, lead to values of the parameters that do not differ significantly. The only relevant exception is $\\tau\\acc$, the typical timescale of accretion onto the MBH, which shows a clear trend to diminish with increasing $z$ as actually expected. It is worth emphasizing that the values found for all the parameters are very reasonable. In particular, those of $\\tau\\acc$, ranging from $10^6$ to $2\\times 10^8$ yr according to $z$, are consistent with the most recent estimates (\\citealt{yu02, schir03}). The ratio $\\epsilon\\gal\\equiv M\\BH/M\\gal$ is found to be $\\simeq 2.5 \\times 10^{-3}$, in agreement with observations \\citep{fermer00,gebhardt00}. The fraction $\\epsilon\\g$ of hot gas that ends up in the central MBH after a MM in the supply-limited accretion regime is of a few percent. The value $\\sim 2$ of the parameter $k$ determining the accretion rate in the Eddington-limited regime is just what one would expect from \\eq (\\ref{ldet}) for a quasar of luminosity $L_8$, representative of a MBH of $10^8$ M$_{\\odot}$, and a typical accretion rate of $\\dot M_8$ during a duty cycle of $10^8$ yr. On the other hand, according to \\cite{cavitt00}, the transition between the Eddington-limited and the supply-limited accretion regimes is expected to occur at the typical formation of groups, at $z$ equal to 2.5$-$3.0, where MMs of haloes of the galactic scale become rarer and the accretion onto MBHs progressively declines. In our model, such a transition takes place when the gas fraction in baryons becomes so low that the product $\\epsilon\\g M\\g$ falls below $k M\\BH$, thus becoming the limiting factor. This occurs at the epoch where $M\\gal/M\\g \\simeq \\epsilon\\g/(k\\epsilon\\gal)$, with the latter ratio taking a value of order unity, which is satisfied at $z\\refe \\sim 2$ not far from the redshifts preferred by \\citet{cavitt00}. The hosting halo masses range from $\\sim 10^{11}$ M$_{\\odot}$ to $\\sim 10^{14}$ M$_{\\odot}$, implying that quasars would be found within central galaxies of haloes from the galactic scale to the scale of galaxy groups, the structures that are currently believed to be the harbouring environments of the brightest quasars (\\citealt{jager01}). Finally, the inferred fraction of observed quasars per halo MM is found to be $\\nu\\simeq$ 5--10\\%, again pretty close to the expectations.  All these results, \\ie the good behaviour of the predicted QLFs at $z>1$ and the reasonable best-fitting values of the parameters, give strong support to the proposed connection between halo MMs and AGN lightning through the triggered simultaneous feeding of the spheroid of the main galaxy and its central MBH. The fact that our model does not match the observed evolution of QLFs at $z<1$ is also not unexpected in this scheme because of the failure for late times of that central hypothesis. The two complementary mechanisms connecting halo MMs and the reactivation of AGN imply the parallel growth of the host galaxy (by cooling flows and the merger of the main galaxies of the progenitor haloes). Consequently, they can be expected to operate as far as haloes are typically of the galactic scale, {\\it when galaxies are growing}. When haloes reach the scale of galaxy groups, galaxies within them no longer grow in size, but just in number. Therefore, none of the above mentioned processes may then be effective. At such late epochs, there may still be radial flows of gas into MBHs lying at the centre of galaxies, but they should be triggered by other mechanisms ---such as internal instabilities of galaxies or tidal interactions among galaxies within groups and clusters--- uncorrelated with halo MMs.  These results also indicate the necessity of dealing more accurately with the physics of baryons, that is, with the processes governing the evolution of the hot gas trapped in haloes, as well as with the growth of galaxies and their interactions. In other words, a QLF model spanning from high to small redshifts should necessarily be coupled to a detailed model of galaxy formation and evolution as in the recent work by \\citet{menci03}. There are other simplifications in our model that might be also improved in future works, although they do not seem to play a crucial role on our main conclusions. There is, for instance, no consideration of the detailed physics at the vicinity of the MBH and our radiation model does not account for relativistic effects. Detailed study of the physics around the MBH could provide, for example, a natural explanation to the decrease of the number of bright quasars (\\citealt{nitta99}).  Likewise, the use of an isotropic model to describe highly anisotropic radiation fields around quasars is likely to introduce different kinds of biases. The first one, due to the requirement that the line-of-sight lies in the emission cone, is of purely statistical nature and should be reasonably accounted for by the normalisation of the MM rate. A second source of bias arises from the relation between the bolometric luminosity and the MBH mass, through the Eddington efficiency, which assumes isotropy of the emission. This approximation, which cannot be corrected by a shift of the luminosity function along the magnitude axis, results in erroneous masses associated with the observed luminosities and, therefore, in erroneous merger rates.  Finally, one must not forget that the comparison with observational data usually takes into account only statistical errors (very small in such large samples as the 2QZ). Systematic effects, however, like the variations of the ability of redshift determination depending on the emission lines present inside the observable wavelenght range, are likely to have a major contribution (\\citealt{croom01a}).  \\vspace{0.75cm} \\par\\noindent {\\bf ACKNOWLEDGMENTS} \\par  \\noindent We would like to thank the people of the Toulouse Extragalactic Team as well as Hugo Capelato, Gary Mamon, Joseph Tapia and many others for help and stimulating discussions, and Brian Boyle for kindly providing us with the latest QLF of the 2QZ survey. G.M.\\ acknowledges funding from \\'Egide and hospitality from IAG-USP and INPE, while E.S.S.\\ acknowledges funding from the Observatoire Midi-Pyr\\'en\\'ees and the hospitality of its staff."
    },
    "0212/astro-ph0212059_arXiv.txt": {
        "abstract": "We present a set of three-component St\\\"ackel potentials defined by five parameters and designed to model the Milky Way. We review the fundamental constraints that any model of the Milky Way must satisfy, including the most recent ones derived from Hipparcos data, and we study how the parameters of the presented potentials can vary in order to match these constraints. Five different valid potentials are presented and analyzed in detail: they are designed to be confronted with kinematical surveys in the future, by the construction of three-integral analytic distribution functions. ",
        "introduction": "The determination of the mass distribution and dynamical structure of the Milky Way is one of the fundamental tasks of Galactic Astronomy. Even though the Milky Way is a spiral barred galaxy, axisymmetric models are a necessary starting point for perturbation analysis and are thus a prerequisite if one wants to understand the bar on a theoretical basis. Caldwell \\& Ostriker\\ (1981), Rohlfs \\& Kreitschmann\\ (1988) and, recently, Dehnen \\& Binney\\ (1998) fitted axisymmetric mass models of the Milky Way to various measurements of the gravitational force field: they concluded that a wide variety of models can emerge from this fitting process and that the mass distribution of the Galaxy is still ill-determined.  In order to fully exploit kinematical stellar surveys, we should construct dynamical models based on Jeans\\ (1915) theorem. This theorem states that the phase space distribution function of a stellar system in a steady state depends only on three isolating integrals of the motion: numerical experiments\\ (Ollongren 1962; Innanen \\& Papp 1977; Richstone 1982) showed that most orbits in realistic galactic potentials admit three such integrals. The third integral, in addition to the binding energy and the vertical component of the angular momentum, is  not analytic in a general potential: it is possible to define an approximate third integral specific to particular orbital families\\ (de Zeeuw, Evans \\& Schwarzschild 1996; Evans, H\\\"afner \\& de Zeeuw 1997; De Bruyne, Leeuwin \\& Dejonghe 2000) or to foliate phase space with tori on which three action integrals can be defined in a potential that differs from the non-integrable one by only a small amount\\ (Kaasalainen \\& Binney 1994; Binney 2002). Instead, we choose to construct models with an exact analytic third integral by using St\\\"ackel potentials\\ (St\\\"ackel 1890). These potentials were introduced into stellar dynamics by Eddington\\ (1915) and have since been used in a number of papers\\ (e.g. Lynden-Bell 1962; de Zeeuw 1985; Dejonghe 1993; Sevenster, Dejonghe \\& Habing 1995; Durand, Dejonghe \\& Acker 1996; Bienaym\\'e 1999) : in fact, the regularity of typical galactic potentials can be understood in terms of their proximity to St\\\"ackel potentials\\ (e.g. Gerhard 1985).  Our long term goal is to constrain the mass distribution of the Galaxy by using kinematical stellar surveys: we shall choose a St\\\"ackel potential, and use the quadratic programming technique described by Dejonghe\\ (1989) to determine, for the available surveys, distribution functions in the space of the integrals of the motion\\ (see Famaey, Van Caelenberg \\& Dejonghe 2002). Then we shall compare the resulting predictions with the surveys. The potential will be modified in the light of that comparison. Thus the first step is to show that different St\\\"ackel potentials are compatible with the standard constraints for a mass model of the Milky Way.  In this paper, our goal is to show that a wide variety of simple St\\\"ackel potentials can fit most known parameters of the Milky Way (including Hipparcos latest findings). In order to do this, we continue the work of Batsleer \\& Dejonghe\\ (1994, hereafter BD) who presented a set of simple St\\\"ackel potentials with two mass components (halo and disc) and a flat rotation curve, that we generalize by adding a thick disc to them since its existence as a separate stellar component is now well documented\\ (Ojha et al. 1994; Chen et al. 2001). These new potentials are described by five parameters and we will show that many different combinations of these parameters are consistent with fundamental constraints for a mass model of the Milky Way.  In the next section, we review the observational constraints that any mass model of the Milky Way must satisfy. In section 3, we present the mathematical form of the set of St\\\"ackel potentials studied in this paper, and we choose selection criteria in the light of the constraints reviewed in section 2. In section 4, we examine the consequences of those selection criteria and we present five potentials differing in terms of form and features, and that are plausible for the Milky Way. For the conclusions, we refer to section 5. ",
        "conclusions": "In this paper, we have shown that some different simple St\\\"ackel potentials can fit most known parameters of the Milky Way (especially Hipparcos latest findings). First, we have reviewed the galactic fundamental parameters that any Milky Way potential must match and that investigations following the Hipparcos mission have readjusted. Then, we have generalized the two-component BD potentials by adding a thick disc, and we have studied how the parameters can vary in order to satisfy selection criteria based on the latest observational constraints. We have shown that the presence of a thick disc allows more flexibility in the form of the potentials, especially for the shape of the halo and we have selected five different valid potentials listed in Table 5. It should be noted that, in fact, there could be two different thick discs in the Galaxy, a very thick one (Chiba \\& Beers 2000; Gilmore, Wyse \\& Norris 2002) and a smaller one (Soubiran, Bienaym\\'e \\& Siebert 2002): in that case, the three-component modelling presented in this paper could be easily extended, but this would imply a growth of parameter space.  A major result of this paper is that, even though St\\\"ackel potentials are negligible in the set of all potentials, many of them are still able to match the most recent estimates for the parameters of the Milky Way, and furthermore very simple ones (superpositions of three Kuzmin-Kutuzov potentials) are sufficient to do this.  The potentials defined in this paper have already been used in Famaey et al.\\ (2002). We plan to further validate each of the five proposed potentials by confronting them with kinematical surveys. This we shall do by constructing three-integral analytic distribution functions in those potentials, using a quadratic programming technique\\ (Dejonghe 1989; for an overview see Dejonghe et al. 2001)."
    },
    "0212/astro-ph0212573_arXiv.txt": {
        "abstract": "Current cosmological observations show a strong signature of the existence of a dark energy component with negative pressure.  The most obvious candidate for this dark energy is the cosmological constant (with the equation of state $w_X = p/\\rho = -1$), which, however, raises several theoretical difficulties. This has led to models for  dark energy component which evolves with time. We discuss certain questions related to the determination of the nature of dark energy component from observations of high redshift supernova. The main results of our analysis are: (i) Even if the precise value of $w_X$ is known from observations, it is {\\em not} possible to determine the nature of the unknown dark energy source using only kinematical and geometrical measurements. We have given explicit examples to show that different types of sources can give rise to a given $w_X$. (ii) Although the full data set of supernova observations (which are currently available) strongly rule out models without dark energy, the high ($z > 0.25$) and low ($z < 0.25$) redshift data sets, individually, admit decelerating models with zero dark energy. Any possible evolution in the absolute magnitude of the supernovae, if detected, might allow the decelerating models to be consistent with the data. (iii) We have introduced two parameters, which can be obtained entirely from theory, to study the sensitivity of the luminosity distance on $w_X$. Using these two parameters, we have argued that although one can  determine the present value of $w_X$ accurately from the data, one cannot constrain the evolution of $w_X$. ",
        "introduction": "One of the equations governing the dynamics of a Friedman universe, $(\\ddot a/a) = - (4\\pi G/3)(\\rho+3p)$, implies that the universe will accelerate ($\\ddot a >0$) if $(\\rho + 3p) <0$. The analysis of high-redshift supernova data \\cite{rfc++98,pag++99,riess00} seems to suggest that $\\ddot a$ is indeed positive thereby requiring $(\\rho + 3p) <0$. This condition requires at least $\\rho$ or $p$ to be negative. Cosmologists have not yet become desperate enough to suggest  $\\rho<0$; but there was remarkably low resistance in the community to accepting the existence  of a constituent with  $w \\equiv (p/\\rho)<-(1/3)$. This was, to a great extent, facilitated by the fact that cosmologists have long since toyed with the existence of a non-zero cosmological constant with the stress-tensor $T^i_k = \\rho_{\\Lambda} \\delta^i_k$ corresponding to $w_{\\Lambda}=-1$. In fact,  much before the arrival of supernova data on the scene, there were indications that the universe actually may be populated by a energy density which is very smoothly distributed over large scales. A very clear argument to this effect was given based on the analysis of APM data in 1990 \\cite{esm90}. Later, the first analysis of the COBE data in 1992 indicated that the shape of the standard cold dark matter (SCDM) spectrum needs to be modified and the existence of a smoother energy distribution, like the cosmological constant will be required \\cite{pn92,ebw92}. An analysis of a host of observations available in 1996 was used to suggest that the data supports the existence of a non-zero cosmological constant \\cite{bpn96}. Given this background, it was not surprising that the cosmologists were not too shocked by the supernova data which became available from 1998 onwards. Currently, there are also other observations, like those of the age of the universe, gravitational lensing surveys etc., which strongly suggest the presence of a positive cosmological constant [see \\citeN{padmanabhan02c} for a comprehensive review and references].  There are, however, well known deep theoretical problems with the existence of a non-zero cosmological constant $\\Lambda$ with a magnitude of about $\\Lambda (G\\hbar/c^3) \\approx 10^{-123}$. [These are well documented in the literature and will not be discussed here; for a recent review see \\citeN{padmanabhan02c}.] This has prompted  a host of activity in which one looks for a dark energy component in the universe with $w<-(1/3)$ which is different from cosmological constant [see \\citeN{padmanabhan02c} and references cited therein]. If $w_X(a)$ denotes the time dependent equation of state parameter for an unknown dark energy component in the universe, then the equation of motion $\\de (\\rho_X a^3) =  w_X \\rho_X \\de a^3$ integrates to give \\begin{equation} \\rho_X(a) = \\rho_X(a_0) \\exp \\left\\{- 3 \\int_{a_0}^a \\frac{\\de a'}{a'} [1 + w_X(a')] \\right\\} \\label{evolenergy} \\end{equation} If $w_X$ is not identically equal to $-1$, then the dark energy density will evolve with time (even if $w_X$ is a constant). This suggests the possibility that the cosmological ``constant'' can be time dependent thereby alleviating some of the difficulties which arises in the case of $w_X=-1$. Several such models (usually based on scalar fields) have been constructed in the literature, all of which, generically have a time dependent $w_X(a)$. The introduction of such an ad-hoc, time dependent \\emph{function} into cosmology, of course, takes away much of the credibility; nevertheless, this issue needs to be settled ultimately by observations and not by aesthetics. (If the observations forces us into a solution which is unaesthetic by current standards, we will either conveniently change the standards or live with it!)  In this paper, we discuss certain questions related to the determination of the nature of dark energy component from observations related to supernova. This issue has been studied in great detail by several groups and our contribution in this paper will border on pedagogy and will present a particular point of view. In the next section we briefly recall and stress some inherent theoretical degeneracies in these analysis. In Section \\ref{sndata} we reanalyze the currently available supernova data in order to focus attention on some key elements. (These results exist in alternate forms in published literature but we believe our analysis brings out some features clearly.) Based on the lessons learnt in this section, we carry out a similar analysis for possible future supernova observations and point out what can and cannot be achieved. We have intentionally kept the data analysis at a fairly simple level (indicated in the title by `a theoretician's analysis'). We subscribe to the point of view that any result which cannot be revealed by a simple analysis of data, but arises through a more complex statistical procedure, is inherently suspect and a conclusion as important as the existence of dark energy with $p<0$ should pass such a test. ",
        "conclusions": "The supernova data on the whole, rules out non-accelerating models with very high significance level, as noted and stressed by various authors \\cite{rfc++98,pag++99,riess00}. It is thus more or less conclusive that we have to look for some form of matter which has negative pressure. However, it is interesting to note that if we divide the data set into high and low redshift subsets, neither of the subsets are able to rule out the decelerating models -- it is only the interplay between the high and low redshift data which implies the result quoted above. This analysis indirectly stresses the importance of any evolutionary effects. Any possible evolution in the absolute magnitude of the supernovae, if detected, might allow the decelerating models to be consistent with the data. However, our understanding of the complicated physical processes in the supernovae is definitely not well enough to discuss the possibility of detecting supernova evolution in near future. Although the intrinsic evolution of the SN is an interesting problem, addressing these issues is beyond the scope of this paper.  The supernova data available today is not enough to fix the parameters $\\Omega_m$ and $\\Omega_{\\Lambda}$ independently. Instead, a particular combination of the two parameters is chosen out by the data. This degeneracy has to do with the dependence of the luminosity distance on these parameters. Since supernova observations essentially measure the luminosity distance $Q(z)$, the accuracy in the determination of various parameters from supernova observations will crucially depend on how sensitive $Q$ is to the changes in those parameters. This analysis can be done entirely from theory [using Fisher information matrix; see for example \\citeN{efstathiou99}] and, in principle, can provide nice handle on which combinations of parameters are expected to be constrained from the supernova data.  The key issue regarding dark energy is whether it is a constant or whether it is evolving with time. In particular, one is interested to determine the evolution of the equation of state, $w_X$ of the dark energy component. To address this question, we have taken a simple phenomenological model for $w_X$ having the form $w_X(a) = w_0 - w_1 (a - 1)$. The sensitivity of the luminosity distance (and hence, the supernova data) on $w_0$ and $w_1$ can be measured by the two parameters ($A$ and $B$) introduced in the paper, which are essentially  the fractional changes in $Q(z)$ for unit change in $w_0$ and $w_1$, respectively. The parameters $A$ and $B$ can be obtained entirely from theory and the conclusions drawn from  them are valid for current as well as future supernova observations. We find that $Q(z)$ is quite sensitive to $w_0$ -- hence, one can constrain the current value of $w_X$ quite well. However, $Q(z)$ is comparatively insensitive to $w_1$, thus determining the evolution of $w_X$ will be a difficult task. The situation is further worsened when we take the uncertainties in $\\Omega_m$ into account. Our results are consistent with other statistical analyses done by \\citeN{astier00}, \\citeN{mbs01}, \\citeN{wa01}, \\citeN{mbms02}, \\citeN{wa02}. They have shown that it is possible to constrain the present value of $w_X$ with future missions like ``SuperNovae Acceleration Probe'' (SNAP); however, constraining the evolution of $w_X$ will not be easy using SNAP, particularly when $\\Omega_m$ is not known to a high accuracy."
    },
    "0212/astro-ph0212090_arXiv.txt": {
        "abstract": "We have developed a timing analysis method to determine the distances of variable galactic X-ray sources based on the method advanced by Tr\\\"{u}mper and Sch\\\"{o}nfelder in 1973. The light-curve of the halo produced by the scattering of X-rays off the interstellar dust is delayed and smeared by the dust grains. This method utilizes the differences between the power density spectra of the point source and the halo. We present the details of this method and our first applications of this method to the Chandra data of X-ray binary Cyg X-3. ",
        "introduction": "When the high energy photons emitted by the X-ray source arrive to the earth, they have been scattered by the interstellar dust grains. Therefore many galactic X-ray sources are observed to have faint and diffuse halos. This effect and its possible application was first discussed by Overbeck (1965), and the X-ray halo was first detected with Einstein Observatory (Rolf 1983, Catura 1983). The investigation of the shape and strength of the haloes was done by several authors (e.g. Mauche \\& Gorenstein 1986, Predehl \\& Schmitt 1995) based on the data of Einstein and ROSAT.  As the photons in the halo travel through longer paths than the photons in the point source, any intensity variation in the the source induces a delayed signal in the halo. It has been suggested by Tr\\\"{u}mper and Sch\\\"{o}nfelder (1973, hereafter TS1973) that this effect can be uesd to determine the distance of variable galactic X-ray sources. Although the method is very simple, the first successful application was realized by Predehl (2000) for a X-ray binary Cyg X-3 almost 30 years later. They obtained the distance of $9^{+4}_{-2}$ kpc to the source. The difficulty of the method is that it is hard to distinguish the halo from the image induced by the point spread function of the instrument until the Chandra X-ray Observatory provides us the high resolution image of X-ray sources.  The shape and strength of the halo depend on the distance to the X-ray source, the geometric distribution and the physical properties of the interstellar dust grains (Hayakawa (1970), Mathis \\& Lee (1991), and Predehl \\& Klose (1996)). ",
        "conclusions": ""
    },
    "0212/astro-ph0212435_arXiv.txt": {
        "abstract": "We present BVRI CCD aperture polarization and near-infrared photometry of the proto-planetary nebula Hen 3-1475. Its intrinsic polarization is high and shows a strong spectral dependence. The position angles in all bands are perpendicular to the axis of the observed bipolar structure. A Monte Carlo code is used to model the intrinsic polarization of \\hhe. Using disk dimensions and other constraints suggested by previous works, we are able to reproduce the observations with an optically thick disk composed by grains with a power-law size distribution ranging from 0.06 to 0.22 $\\mu$m. We also reliably estimate the foreground polarization from hundreds of stars contained in the CCD images. It is parallel to the intrinsic polarization of Hen 3-1475. Possible implications of this result are discussed. From IR observations, we estimate a interstellar reddening, A$_V$, of about 3.2. ",
        "introduction": "\\he (IRAS 17423-1755) is identified as a proto-planetary nebula (PPN), i.e., an object between the end of the asymptotic giant branch (AGB) and the planetary nebula (PN) phases. The classification and physical properties of Hen~3-1475 are supported by a broad base of observational data: radio emission  (Knapp et al. 1995); OH masers (te Lintel Hekkert 1991; Zijlstra et al. 2001); optical imaging and spectroscopy (Riera et al. 1995; Bobrowski et al. 1995; Borkowski, Blondin, \\& Harrington 1997; Borkowski \\& Harrington 2001). Probably the strongest evidence in favor of its evolutionary stage is the high N abundance, typical of evolved objects, specifically of Type I PN (Riera et al. 1995). The radio continuum indicates the presence of a compact ionized region  surrounding the central B star with an estimated dynamical age of only 15 years (Knapp et al. 1995. Recently, Borkowski \\& Harrington (2001) have obtained images and detailed spectroscopic observations and, among other results, they have  derived a distance of 8~kpc, which implies a very high luminosity (25\\,000 L$_\\odot$) for a AGB star at the Galactic Bulge. Analyzing H$\\alpha$ spectra, S\\'anchez Contreras \\& Sahai (2001) could identify two winds, one of which is probably associated with the post-AGB wind before its interaction with the previous AGB wind.  High-resolution optical images of Hen~3-1475 show a very complex bipolar morphology (Borkowski et al. 1997; Bobrowski et al. 1995; Riera et al. 1995). The central region is dense, probably a dust disk, which may be responsible for the OH maser emission as well as for the CaII and FeII lines observed in its near-infrared spectrum. This structure has an angular size of about 2\\arcsec. The high resolution image obtained by Borkowski et al. (1997) shows also dark patches probably associated with denser clumps of material in the disk. A complex structure of highly collimated optical jets with pairs of point-symmetrical shock-excited knots produces a characteristic S-shaped sequence, which is about 15\\arcsec\\ in size oriented along PA 135$^o$.  In the rapid evolutionary phase of PPN, dramatic changes in geometrical and physical properties of its central star and circumstellar material take place. An example is the development of non-symmetric geometries (eg., bipolar nebulae) in post-AGB objects. The interaction between an equatorial enhanced AGB wind and that from an early PN phase was suggested as a mechanism to produce the variety of geometries seen in PPNe and PNe (Kahn \\& West 1985). The status of this and other models is reviewed by Frank (2000).  This enhanced equatorial wind - which we will call disk - can also play a role in collimating the jets seen in many objects. Recently, Frank, Lery, \\& Blackman (2002) have shown that a magnetic-centrifugal launching model can explain winds with velocities as high as those observed in \\hhe.  Therefore, a good representation of the disk properties is crucial to understand the changes occurring during the PPN phase.  Polarimetry is an useful tool in the study of PNe and PPNe because it is sensitive to the circumstellar geometry and composition. The polarization measured in PPNe is usually high due to the scattering in asymmetric circumstellar environments. Johnson \\& Jones (1991) have shown that, from red giants to PNe, the evolutionary stage of PPN is the one in which the polarization tends to be highest and that the transition from a spherically symmetric envelope (typical of red giants) to a non-spherical envelope seems to occur in early stages of the AGB phase. The wavelength dependence of the polarization in PPNe has various shapes and some objects present polarization-angle rotation (Trammell, Dinerstein, \\& Goodrich 1994). Many extended objects were observed using imaging polarimetry (e.g., Scarrot \\& Scarrot 1995; Gledhill et al. 2001). Far from the central source, the light is highly polarized and exhibits the centro-symmetric pattern typical of reflection nebulae. However, the polarization near the central object tends to be parallel to the optically thick disk originated in the AGB phase. This pattern is also observed in young stellar objects with bipolar geometries (see references in Bastien \\& Menard 1988 - Table 1).  In spite of the importance of grain size in the study of the temperature profile and, consequently, of the infrared (IR) luminosity (Stasi\\'nska \\& Szczerba 1999, 2001), few works have focused on the determination of the grain size distribution in PPNe and PNe. An example of model that considers the grain sizes as free parameters is found in Meixner et al. (2002). That work, however, modeled the observed spectral energy distribution (SED) alone. Recent results (Carciofi, Bjorkman, \\& Magalh\\~aes 2002; Li \\& Lunine 2002) show that the SED alone does not constrain the grain sizes. The spectral dependence of the polarization helps to restrict the models to the dust size distribution (Carciofi, Magalh\\~aes, \\& Bjorkman 2002). Heckert \\& Smith (1988) modeled the polarization of OH 0739-14 considering single scattering in an envelope composed by grains of a single size. Johnson \\& Jones (1991) have modeled the observed polarization of some PPNe, but they used a fixed MRN size distribution ranging from 0.005 to 0.10 $\\mu$m. To our knowledge, there is no previous attempt to describe the polarization data of a PPN with a free distribution of grain sizes and multiple scattering.  In this work, we present data on broad band CCD aperture polarimetry as well as near-infrared (NIR) photometry of Hen~3-1475 (Section \\ref{sec_obs}). In Section \\ref{sec_ir} the NIR photometry is analyzed in order to provide information on the interstellar (IS) and intrinsic reddening. The polarimetric observations are discussed and modeled in Section \\ref{sec_polar} providing some information about the Hen~3-1475 dusty envelope. In the last section, the main conclusions are listed. ",
        "conclusions": "In this work we analyze BVRI polarization and IR photometry of the proto-planetary nebula Hen 3-1475. The main results are summarized below:  \\begin{itemize}  \\item the optical broad band fluxes of \\he are highly polarized. A large fraction of this polarization may be produced in the circumstellar environment of this object;  \\item the BVRI polarization position angle is perpendicular to the axis of the bipolar structure seen in optical images. As the aperture used to measure the polarization is small and therefore does not include the jets, that implies that the central disk must be optically thick in all the bands observed or the dust must be composed of a highly absorptive material;  \\item the wavelength dependence of the polarization of \\he was modeled using a Monte Carlo radiative transfer code. We could find a good fit using an index -3.0 for the power-law grain-size distribution, whose limits are 0.063 and 0.22 $\\mu$m. In spite of the claim for an oxygen chemistry envelope (Knapp et al. 2000), the optical constants of amorphous carbon produce a better fit to the optical polarization than those of silicates. However, the simple model used here does not allow to discard silicate dust in the \\he circumstellar envelope;  \\item our CCD images provide us with polarimetric data of a large number of stars in the \\he field. So we could reliably estimate the foreground polarization in this direction. Curiously, the foreground polarization has the same direction of the intrinsic polarization of \\hhe. This might be interpreted as an interplay between the IS magnetic field and the \\he geometry;  \\item the IR colors of \\he suggests an interstellar reddening in the V band of around 3.2 mag in the line of sight. The circumstellar extinction was estimated as A(V) $\\approx$ 9.0 mag.  \\end{itemize}"
    },
    "0212/astro-ph0212329_arXiv.txt": {
        "abstract": "We perform Monte Carlo simulations of diffusive shock acceleration at highly relativistic oblique shock waves. High upstream flow Lorentz gamma factors ($\\Gamma$) are used, which are relevant to models of ultra relativistic particle shock acceleration in Active Galactic Nuclei (AGN) central engines and relativistic jets and Gamma Ray Burst (GRB) fireballs. We investigate numerically the acceleration properties in the relativistic and ultra relativistic flow regime ($\\Gamma \\sim 10-10^{3}$), such as angular distribution, acceleration time constant, particle energy gain versus number of crossings and spectral shapes. We perform calculations for sub-luminal and super-luminal shocks. For the first case, the dependence on whether or not the scattering is pitch angle diffusion or large angle scattering is studied. The large angle model exhibits a distinctive structure in the basic power-law spectrum which is not nearly so obvious for small angle scattering. However, both models yield significant 'speed-up' or faster acceleration rates when compared with the conventional, non-relativistic expression for the time constant, or alternatively with the time scale $r_{g}/c$ where $r_{g}$ is Larmor radius. The $\\Gamma^{2}$ energization for the first crossing cycle and the significantly large energy gain for subsequent crossings as well as the high 'speed up' factors found, are important in supporting the Vietri and Waxman work on GRB ultra-high energy cosmic ray, neutrino, and gamma-ray output. Secondly, for super-luminal shocks, we calculate the energy gain for a number of different inclinations and the spectral shapes of the accelerated particles are given. In this investigation we consider only large angle scattering, partly because of computational time limitations and partly because this model provides the most favourable situation for acceleration. We use high gamma flows with Lorentz factors in the range 10-40, which are relevant to AGN accretion disks and jet ultra relativistic shock configurations. We closely follow the particle's trajectory along the magnetic field lines during shock crossings where the equivalent of a guiding centre approximation is inappropriate, constantly  measuring its phase space co-ordinates in the fluid frames where {\\bf E}=0. We find that a super-luminal 'shock drift' mechanism is less efficient in accelerating particles to the highest energies observed, compared to the first order Fermi acceleration applying in the sub-luminal case, suggesting that the former cannot stand as a sole acceleration mechanism for the ultra-high energy cosmic rays observed. ",
        "introduction": "Detection of $10^{6}-10^{20}$eV nuclei within the solar system, implying their presence through much of the observable universe, has stimulated the development of the diffusive shock acceleration model, whereby particles are repeatedly accelerated in multiple crossings of a shock interface, due to collisions with upstream and downstream magnetic scattering centres. The pioneering work of the late 70s (Krymsky, 1977; Bell, 1978a,b; Axford et al., 1978; Blandford and Ostriker, 1978) established the basic mechanism of particle diffusive acceleration in non-relativistic flows.  However, since then, extension of the model to oblique shocks, relativistic plasma flows and the incorporation of  non-linear effects, have rendered the emergence of a consensus physical picture more difficult. At oblique shock fronts the motion of the particles is more complicated, compared to a  parallel shock configuration, as they may either be transmitted or reflected by the shock surface. Apart from the diffusive acceleration mechanism, an oblique shock could also accelerate particles in the \\textit{absence} of scattering (Armstrong and Decker, 1979), and this mechanism is called 'shock drift' acceleration. Provided that the motion of the particles is diffusive in the oblique shock configurations it has been shown, for an analytical non-relativistic flow approach, that the same result as in parallel shocks, connects the power-law index of accelerated particles with the compression ratio and the acceleration rate with the diffusion coefficients in the shock normal direction (e.g. Axford et al., 1978; Bell, 1978a; Drury, 1983).\\\\ Also, following energy gain in the shock frame on shock crossing via the drift approach, is equivalent to following the gain via Lorentz transformations from the shock frame into an {\\bf E}=0 frame (defined by de Hoffmann and Teller, 1950) at the shock, and then back into the shock frame. On the other hand, the theory of oblique shocks considers features that are not involved in the parallel shock case, in particular the ability of the shock to reflect particles meeting a weaker to stronger field transition and the dependence of the acceleration time on the field-shock normal angle (e.g. Jokipii, 1982). Oblique shocks can be classified into sub-luminal and super-luminal. The difference between these two categories is that, in the sub-luminal case, it is possible to find a relativistic transformation to the frame of reference (de Hoffmann-Teller frame), in which the shock front is stationary and the electric field is zero ({\\bf E}=0) in both upstream and downstream regions. On the other hand, super-luminal shock fronts do not admit a transformation to such a frame, as they correspond to shock fronts in which the point of intersection of the front with a magnetic field line moves with a speed greater than $c$. Kirk and Heavens (1989) employed a numerical solution to the transport equation (in the test particle limit) and found spectral flattening at relativistic speeds in inclined shocks as well as an increased anisotropy both in fluid frames and the {\\bf E}=0 frame. These results were confirmed by Monte Carlo simulations (Ostrowski, 1991 and Lieu et al., 1994). The former author investigated the gyromotion across the shock to check the region of applicability of magnetic moment conservation in the {\\bf E}=0 frame at the shock while, the later authors simply relied on this assumption. Baring et al. (1994) noticed 'bumps' in the accelerated spectra just upstream of the shock in the non-relativistic oblique case. The reason a Monte Carlo approach became favoured, lies in the large anisotropies and gradients which develop around the shock at relativistic flow speeds. These render direct solution of the transport equation difficult, because the Boltzmann equation must be applied to a fine-grained bundle of particle trajectories in phase space and the diffusion approximation, employing a simple spatial diffusion tensor is inadequate (e.g. Axford, 1981).\\\\ All the work mentioned so far is sub-luminal, but Bednarz and Ostrowski (1998) used a small-angle scattering model to include  super-luminal situations with field-shock normal angles, $\\psi \\leq \\pi/2 $ and flow $\\Gamma \\leq 243$ and also a simulation of varying amounts of cross-field diffusion. These authors found spectral indices which could be very large at intermediate $\\Gamma$ values with small $K_{\\perp}$, but at high $\\Gamma$, all results appeared to tend towards a differential exponent of -2.2. \\\\ Earlier discussion of an extreme example of a super-luminal case, the pulsar wind shock acceleration scheme for the termination shock of an isolated pulsar magnetosphere where the $\\Gamma \\sim 10^{4}$ polar wind meets the slowly expanding nebula  envelope, has emphasized the apparent difficulty in achieving multiple shock crossings of accelerated particles. This can be seen if we remember that for super-luminal shocks, a transformation to a frame with {\\bf B} parallel to the shock plane is possible (Hudson, 1965) so in the absence of scattering close to the shock, particles are swept at an angle to the shock normal into the shock, shock drift and then exit downstream with no return possible. Begelman and Kirk (1990) note that only shock compression of the upstream population results in the low turbulence case. However, Lucek and Bell (1994) find special trajectories if the field is ordered but not uniform, allowing energy increases by factors of the order of hundreds per crossing, which coupled with particle motion in the electromagnetic field of the nebula can render isolated pulsars accelerators of high energy cosmic rays. Gallant and Arons (1994) provide a rare example of a hybrid model of relativistic particle acceleration, applied to the Crab Nebula, using a fluid approach for an $e^{\\pm}$ pair plasma together with a kinematic treatment of ions to compute the synchrotron spectrum. \\\\ Acceleration time under non-relativistic theory clearly depends on shock obliquity through variation of $K_{n}$, the diffusion coefficient along the shock normal, with $\\psi$, where $K_{n}=K_{||}cos^{2}\\psi+K_{\\perp}sin^{2}\\psi$. $K_{||},K_{\\perp}$ are respectively the diffusion coefficients parallel and perpendicular to the field. The acceleration time is given by,  \\begin{equation} \\tau=\\frac{3}{V_{1}-V_{2}}\\left(\\frac{K_{n,1}}{V_{1}}+\\frac{K_{n,2}}{V_{2}}\\right) \\end{equation}  Here the suffixes '1' and '2' refer to upstream and downstream, and $V$ is plasma flow velocity. Because this equation (1) is clearly an inadequate measure of acceleration time in the highly relativistic case, we note it may also be written as $\\tau$ lying in the range $(4\\rightarrow 8/3)(nr_{g}/c)cos^{2}\\psi$ for a parallel mean free path $\\lambda_{\\|}=nr_{g}$, a shock frame compression ratio of 4 and a downstream reduction in $\\lambda_{\\|}$ by a factor 4. Hence while it will be convenient to measure any 'speed-up' of acceleration relative to equation (1), we may equally regard the time to move over the relativistic Larmor radius as the unit of acceleration time for comparison purposes. Quenby and Lieu (1989) and Ellison et al. (1990) employed Monte Carlo guiding centre approximation  calculations to show that relativistic parallel shock flows with $\\Gamma\\sim 10$ produced a 'speed-up' of acceleration by factors $\\sim3\\rightarrow10$ when compared with equation (1) for $\\tau$. The former authors argued for large angle scattering while the latter found essentially the same result with either large or small angle scattering. Lieu et al. (1994) showed the speed-up as a function of flow velocity was similar for inclined, sub-luminal shocks, provided the results are plotted  against flow seen in the {\\bf E}=0  frame. Quenby and Drolias (1995) using an inclined shock model which included the expected shock structure due to cosmic ray pressure back reaction, found some speed-up still occurred. Bednarz and Ostrowski (1996) employ a Monte Carlo computation scheme to investigate acceleration times for shock speeds up to 0.9$c$ for the inclined case, with both  small angle and large angle scattering and allowing for cross field motion of guiding centres. Speed-up is noticed especially at high shock speed and turbulence although, cross-field diffusion eventually reverses the correlation  with turbulence level. Power law spectra only occur in the super-luminal case for isotropic scattering. Acceleration time scales, shorter than the upstream gyroperiod have been obtained.\\\\ In a companion work, Meli and Quenby (2002a), have studied parallel shock acceleration for flow $\\Gamma \\leq 10^{3}$,  for both small and large angle scattering models using the guiding centre Monte Carlo approach. They confirm a $\\Gamma^{2}$ energy enhancement for the first shock crossing cycle, first emphasized by Quenby and Lieu (1989,) and they find significant energy gains subsequently. Also, the structure in the accelerated particle spectrum at the shock, indicated in previous lower $\\Gamma$ simulations, became enhanced for large angle scattering, although not in the small angle case, rendering the idea of a simple power law output spectrum for photon production uncertain and they confirm the existence of the speed-up effect. It is the purpose of this work to extend the parallel shock studies of Meli and Quenby (2002a) to provide answers concerning the acceleration time scales, particle spectral shape at the source, angular distribution and energy gain per cycle of the accelerated particles for different shock inclinations using very high gamma flows. We consider both sub-luminal shocks (where an {\\bf E}=0 frame can be found) and super-luminal ones (where {\\bf E}=0 is impossible) and the helical particle trajectories are followed  in the appropriate frames to determine the characteristics of the shock crossing. Streaming instabilities, drift waves and Rayleigh-Taylor like instabilities all can increase scattering around the shock and contribute to the back reaction of the relativistic gas, but these are not enough to destroy the basic relativistic beaming in flows, where $V \\rightarrow c$. Hence a Monte Carlo approach is employed. The speed-up for a true discontinuity in the test-particle strong scattering limit, is due to a $\\sim \\Gamma ^{2}$ change in energy in one complete shock crossing cycle for shock $\\Gamma$, and this effect is the basic cause of the deviation of the spectral shape from the non-relativistic result. Our current interest, regarding a real astrophysical scenario, lies in relativistic flows in AGN jets where $\\Gamma\\approx 5-10 $ and in GRB fireballs where a $\\Gamma\\approx 10^{2}-10^{3}$ is estimated. ",
        "conclusions": "Oblique particle shock  acceleration simulations for sub-luminal and super-luminal cases have been presented. High gamma upstream flows have been used ranging from $\\Gamma$=10 up to $\\sim 10^{3}$.  The results found are important in answering questions concerning the efficiency and relativistic effects of first order Fermi acceleration. A major finding of both this work and the companion paper (Meli and Quenby, 2002a), is that in contrast to parallel shocks with small angle scattering, all other types of shocks and scattering models do not yield smooth, power law spectra at the shock in the $\\Gamma>>1$ regime, even though subsequent smoothing during particle escape may arise. Reasons for this will be re-iterated later. However, sub-luminal, inclined relativistic shocks provide spectra more extended in energy and therefore more approximating to a power law distribution than super-luminal shocks, the latter exhibiting a pronounced high energy cutoff feature. The result that up to 3 crossings of the shock can often occur in the adopted large angle scattering model implies that the super-luminal case is not simply 'one-shot' shock drift. However the large gain crossing $2 \\rightarrow 3$ followed by loss to the system can well be described as the energy gain experienced on the last shock drift pass through the discontinuity standing out from the smaller, average diffusive gain.\\\\ Concerning the important 'speed up' effect in the sub-luminal case, this appears to be due to a combination of $\\Gamma^{2}$ or $\\Gamma$ energy gain and the anisotropic angular distributions at the shock front (Ellison et al., 1990) which as we showed, is confined to a very small angle transmission cone. Thus, only a narrow range of pitch angle particles return to the shock and are 'counted' for the acceleration time calculation, as the ones with larger pitch angles are lost downstream. If the particle distribution is made isotropic each side of the shock before re-crossing, the energy gain per cycle is $\\sim\\Gamma_{1}^{2}$ Quenby and Lieu (1989), but a distribution remaining anisotropic will experience less energy gain (Gallant and Achterberg, 1999). Hence some part of the 'speed up' compared with the non-relativistic situation is expected to be connected with the un-allowed for energy enhancement. The anisotropy is of course linked to the 'beaming effect' found also in mild relativistic cases, in accordance with Peacock (1981) and Kirk and Schneider (1988). Lieu and Quenby (1990) and Lieu et al. (1994) note that the distribution of the particle's pitch angle is peaked towards the direction of the  relativistic transformation to the de Hoffmann-Teller frame. In other words this beaming which enhances the number of particles close to the field  line, may provide an additional reason for the decrease in acceleration time as the flow becomes more relativistic for both large angle and pitch angle diffusion. This 'speed-up' resembles a trend found in the parallel shock particle acceleration case. Furthermore, the last arguments could explain the prominent spectral 'plateau' shapes found in many oblique, sub-luminal and parallel shock cases where each shock cycle crossing seems to stand  out \\textit{separately}. Enhanced single cycle acceleration is followed by large flux loss downstream, preventing statistical smoothing into a single power law accelerated spectrum.\\\\ The 'speed up' effect -which has also been observed in the mildly relativistic regime Quenby and Lieu (1989) and Ellison et al. (1990), is important for a complete understanding of particle acceleration to the highest energies. An understanding of this crucial effect is necessary in any explanation for the maximum energy that particles (protons, electrons, iron nuclei) can  achieve in GRB fireballs, AGN hotspots, and pulsar ultra relativistic polar winds. For example results found  in the previous sections dealing with the sub-luminal case for $\\Gamma\\sim 10^{3}$, show only few crossings are needed for the particle to achieve energies up to $10^{19}-10^{20}$eV, although for $\\Gamma\\sim$ 10 more than ten crossings may be needed to gain such high energies. It is the competition between loss time and energy gain time which determines the maximum energy in many situations. Generally speaking the percentage increase of the maximum energy would depend on many factors. One of those is the velocity of the upstream flow. Also the dimensions of the shock could play a vital role in order to define the escape losses in a specific shock acceleration site. In addition another factor in the physical conditions which could alter the cosmic ray spectrum and decrease the spectral index is due to different energy losses occurring in the regime; for example the presence of high or low energy photons (Greisen-Zatsepin-Kuz'min cutoff effect{\\footnote {This is called also the microwave background effect, where the interaction between the cosmic rays and the microwave background results in a cutoff in the energy spectrum of cosmic rays (Greisen 1966; Zatsepin-Kuz'min 1966)}}) or $\\gamma-\\gamma$ flux interactions{\\footnote{If GRB high energy photons come from electron acceleration in the $\\Gamma\\sim 10^{3}$ regime, we might expect some 'structure' in the spectrum at the higher energies on any scattering model.} in GRB will decrease the maximum energy of the particles. In the context of AGN it is well known from a number of interesting papers such as Protheroe and Kazanas  (1983) and Chakrabarti and Moltoni (1993) that a shock in the central engine of the accreting flow ($\\Gamma \\simeq 5-10$) is likely to appear and particle acceleration can consequently take place. Also there is observational evidence that a common feature of many AGN is the formation of one or two jets where shock formations should appear in conjunction with $\\Gamma\\sim 5-10$ flows. Thus, the computed flattening of the spectra -they are all flatter than a -2 power law-  as well as the acceleration 'speed up', should have important consequences when compared with predictions and observational data, regarding for example, the photons produced by emission from electrons which  can be accelerated at the shocks or produced from accelerated protons. GRB seem to be another potential candidate for the acceleration of Ultra High Energy Cosmic Rays (UHECRs), but the issue is debated by many workers.  As we have mentioned Vietri (1995) and Waxman (1995), who have exploited the 'gamma-squared' factor energy gain of the particle, were first to predict theoretically that indeed UHECRs can be produced in GRB where a pre-acceleration of the particles has been implied. Briefly, we note that observationally there is very strong evidence that pulsars (e.g. PSR 1913+16, PSR 2127+11C) are capable of injecting continuously relativistic particles into the  surrounding medium over their lifetime. While this plasma will probably be predominantly in the form of $e^{+}e^{-}$ pairs created in the pulsar magnetosphere, it has been argued that pulsar winds must also contain ions in order to account for the electrical current required in the Crab Nebula (Hoshino et al., 1992; Gallant and Arons, 1994). The termination shock as the $\\Gamma\\sim 10^{5}$ flow meets the nebula is likely to be quasi-perpendicular and hence the super-luminal situation applies for the additional $e^{+}e^{-}$ acceleration. The enhanced production of neutrinos in astrophysical sites such as in AGN and GRB is another consequence of the simulation findings of the 'speed up' effect and 'gamma squared' factor.  Mastichiadis and Kirk (1992),  Protheroe and  Szabo (1992) and Vietri (1998) showed that the neutrino flux depends on the maximum energy attained from the primary protons, thus a decrease of the acceleration time constant could produce higher maximum energy limits for the accelerated particles and consequently an increase in the cutoff energy of the neutrino flux will be noted.\\\\   {\\noindent \\bf Acknowledgements}\\\\  We would like to express our thanks to Prof. Drury and the unknown referee for valuable comments and suggestions."
    },
    "0212/hep-ph0212270_arXiv.txt": {
        "abstract": "The first KamLAND results are in a very good agreement with the predictions made on the basis of the solar neutrino data and the LMA realization of the MSW mechanism. We perform a combined analysis of the KamLAND (rate, spectrum) and the solar neutrino data with a free boron neutrino flux $f_B$.  The best fit values of neutrino parameters are $\\Delta m^2 = 7.3 \\cdot 10^{-5}$ eV$^2$, $\\tan^2 \\theta = 0.41$ and $f_B = 1.05$ with the $1\\sigma$ intervals: $\\Delta m^2 = (6.2 - 8.4) \\cdot 10^{-5}$ eV$^2$, $\\tan^2 \\theta = 0.33 - 0.54$.  We find the $3\\sigma$ upper bounds: $\\Delta m^2 < 2.8 \\cdot 10^{-4}$ eV$^2$ and $\\tan^2 \\theta < 0.84$, and the lower bound $\\Delta m^2 > 4 \\cdot 10^{-5}$ eV$^2$.  At 99\\% C.L. the KamLAND spectral result splits the LMA region into two parts with the preferred one at $\\Delta m^2 < 10^{-4}$ eV$^2$.  The higher $\\Delta m^2$ region is accepted at about $2\\sigma$ level.  We show that effects of non-zero 13-mixing, $\\sin^2 \\theta_{13} \\leq 0.04$, are small leading to a slight improvement of the fit in higher $\\Delta m^2$ region.  In the best fit point we predict for SNO: CC/NC $= 0.33 ^{+0.05}_{-0.03}$ and $A_{DN}^{SNO} = 2.8 \\pm 0.8\\%$ (68\\% C.L.), and $A_{DN}^{SNO} < 9 \\%$ at the $3\\sigma$ level. Further improvements in the determination of the oscillation parameters are discussed and implications of the solar neutrino and KamLAND results are considered. ",
        "introduction": "The first KamLAND results~\\cite{kamf} are the last (or almost last) step in resolution of the long-standing solar neutrino problem~\\cite{john}.  ``In the context of two flavor neutrino oscillations with the CPT invariance, the KamLAND results exclude all oscillation solutions of this problem but the `Large mixing angle' solution\"~\\cite{kamf}.  In fact, KamLAND excludes also all {\\it non-oscillation} solutions based on neutrino spin-flip in the magnetic fields of the Sun, on the non-standard neutrino interactions, etc.. More precisely, KamLAND excludes them  as the dominant mechanisms of the solar neutrino conversion.    Soon after the suggestion of the MSW mechanism \\cite{msw}, the ``MSW triangle\"~\\cite{triangle} had been constructed in the $\\Delta m^2 - \\sin^2 2\\theta$ plane which corresponds to 1/4 - 1/3 suppression of the Argon production rate~\\cite{Cl}.  The region around the upper right corner of this triangle, or large $\\sin^2 2\\theta$ part of its horizontal side (``the adiabatic solution\") is what we call now the LMA region.  Kamiokande~\\cite{kamioka}, SAGE~\\cite{sage} and GALLEX~\\cite{gno} have disintegrated the triangle into pieces excluding some parts of oscillation parameter space.  First, the results from Kamiokande (the absence of a strong spectrum distortion as well as the absence of strong Day-Night effect) have chipped off the LMA region. Then Gallium experiments SAGE and GALLEX, confirming LMA, have splitted the diagonal side of the triangle into SMA and LOW.  Apart from that vacuum oscillation solutions (VO) were always present.  For a long time SMA was the favored solution and LMA was considered as an non excluded possibility. Things changed in 1998 when Super-Kamiokande~\\cite{SK} has testified against SMA (the flatness of spectrum and the absence of a peak in the zenith angle distribution of events in the Earth core bin).  Soon after that, the analysis of the Super-Kamiokande data for 708 days of operation allowed us~\\cite{bks} to conclude that solar neutrino results provide hints that the LMA solution could be correct (see also analysis in~\\cite{valle}).  Since that time LMA continuously reinforced its position.   Super-Kamiokande~\\cite{SK} and then SNO~\\cite{sno} have accomplished genesis of the triangle. They have practically excluded SMA, disfavored LOW and VO, and shifted the LMA region to larger $\\Delta m^2$ (small Day-Night asymmetry at Super-Kamiokande) and smaller mixings (small ratio of CC/NC events at SNO). \\\\   By the time of the KamLAND announcement, the solar neutrino data~\\cite{Cl,sage,gno,SK,sno,sno-nc,sno-dn} have definitely selected LMA as the most favorable solution based on neutrino mass and mixing~\\cite{mswlma,us:msw}. The best fit point from the free boron neutrino flux fit~\\cite{us:msw} is \\be \\Delta m^2 = 6.15 \\cdot 10^{-5} {\\rm eV}^2 , ~~~ \\tan^2 \\theta = 0.41,~~~ f_B = 1.05, \\label{bfsol} \\ee where $f_B \\equiv F_B/F_B^{SSM}$ is the boron neutrino flux in the units of the Standard Solar Model predicted flux~\\cite{ssm}.   On basis of the solar neutrino results (and the assumption of the CPT invariance) predictions for the KamLAND experiment have been calculated. A significant suppression of the signal was expected in the case of the LMA solution.  The predicted ratio of the numbers of events with the visible (prompt) energies, $E_p$, above 2.6 MeV  with and without oscillations equals~\\cite{us:msw}: \\be R_{KL}^{LMA} = 0.65 ^{+ 0.09}_{-0.40} ~~~(3\\sigma). \\label{klr} \\ee For other solutions of the solar neutrino problem one expected $R_{KL} = 0.9 - 1$, where the deviation from 1 can be due to the effect of nonzero 1-3 mixing.  In the best fit point (\\ref{bfsol}) the predicted spectrum has (i) a peak at $E_{p} \\approx (3.0 - 3.6)$ MeV, (ii) a suppression of the number of events near the threshold energy $E_{p} \\approx 2.6$ MeV and (iii) a significant suppression of the signal (with respect to the no-oscillation case) at the high energies: $E_{p} > (4 - 5)$ MeV \\cite{us:msw}.  No distortion of the spectrum is expected for the other solutions.\\\\  The first KamLAND results: both the total number of events and spectrum shape~\\cite{kamf}, are in a very good agreement with predictions: \\be R_{KL}^{exp} = 0.611 \\pm 0.094~. \\label{bestkamrates} \\ee The spectral data (although not yet precise) reproduce well the features described above.  As a result, the allowed ``island\" in the $\\Delta m^2 - \\sin^2 2\\theta$ plane with the best fit KamLAND point covers the best fit point from the solar neutrino analysis~\\cite{kamf}.\\\\  In this paper we present our analysis of the KamLAND data as well as combined analysis of the solar neutrinos  and KamLAND results.  The impact of KamLAND on the LMA MSW solution is studied.  We consider implications of the combined (solar neutrinos + KamLAND) analysis and make predictions for the future measurements.  During preparation of this paper, several studies of the first KamLAND results and solar neutrino data have been published~\\cite{after} - \\cite{alia}. Our conclusions are in a good agreement with results of those papers, although there are some differences. Detailed comparison of results will be done later. ",
        "conclusions": "\\noindent 1. The first KamLAND results (rate and spectrum) are in a very good agreement with the predictions based on the LMA MSW solution of the solar neutrino problem.\\\\   \\noindent 2. Our analysis of the KamLAND data reproduces well the results of the collaboration.  The oscillation parameters extracted from the KamLAND data and from the solar neutrino data agree within $1\\sigma$.\\\\   \\noindent 3. We have performed a combined analysis of the solar and KamLAND results.  The main impact of the KamLAND results on the LMA solution can be summarized in the following way. KamLAND  \\begin{itemize}  \\item shifts  the $\\Delta m^2$ to slightly higher values $6.15 \\cdot 10^{-5}$ $\\rightarrow 7.3 \\cdot 10^{-5}$ eV$^2$; (the mixing is practically unchanged: $\\tan^2 \\theta = 0.41$, and this number is rather stable with respect to variations of the analysis;  \\item establishes rather solid lower bound on the $\\Delta m^2$: $\\Delta m^2 > 4 \\cdot 10^{-5}$ eV$^2$ ($3\\sigma$);  \\item disintegrates  the LMA region at 99\\% C.L.  into two parts.  \\end{itemize}  KamLAND further disfavors high values of $\\Delta m^2$: $\\Delta m^2 > 3\\cdot 10^{-4}$ eV$^2$.\\\\  \\noindent 4. Inclusion of the 1-3 mixing   shifts the allowed regions to larger $\\Delta m^2$ and smaller $\\tan^2 \\theta$. $\\chi^2$ slightly increases with $\\sin^2 \\theta_{13}$.\\\\   \\noindent 5. The KamLAND results strengthen the upper bound on the expected value of the day-night asymmetry at SNO: $A_{DN}^{SNO} < 9 \\%$.  We predict about 10\\% turn up of the energy spectrum at SNO in the interval of energies (5 - 8) MeV.  The CC/NC ratio is expected to be CC/NC $\\approx 0.33$, that is, near the present best SNO value.  Future SNO  measurements of the CC/NC ratio and $A_{DN}^{SNO}$ will have further strong impact on the LMA parameter space."
    },
    "0212/astro-ph0212165_arXiv.txt": {
        "abstract": "{We report the detection of short period oscillations in the hot subdwarf B (sdB) star PG 1613+426 from  time-series photometry carried out with the 91-cm Cassegrain telescope of the Catania Astrophysical Observatory. This star, which is brighter than the average of the presently known sdB pulsators, with B = 14.14 mag, has $T_{\\rm eff}=34\\,400\\,{\\rm K}$ and $\\log g = 5.97$, its position is near the hot end of the sdB instability strip, and it is a pulsator with a well observed peak in the power spectrum at $144.18\\pm 0.06\\,\\rm s$. This star seems to be well suited for high precision measurements, which could detect a possible multi-mode pulsation behaviour. ",
        "introduction": "The hot subdwarf B (sdB) stars form a homogeneous group populating an extension of the extreme horizontal branch (EHB) in the $(T_{\\rm eff}- \\log g)$-diagram towards temperatures up to $40\\,000\\,{\\rm K}$. These stars are evolved low-mass $(\\sim 0.5 M_\\odot)$ objects with a He burning core surrounded by a H surface layer which is too thin ($< 4\\%$ by mass) to sustain the H-shell burning. Their origin is still debated, but it seems that they have experienced He-flash phase and a substantial mass loss along the red giant branch. Their further evolution proceeds toward the EHB by crossing the subdwarf O population and eventually entering the white-dwarf graveyard (Maxted et al. 2001, Heber et al. 2002).  The recent discovery that several of them are multimode pulsators has triggered both observational and theoretical efforts for studying their characteristics and pulsating mechanisms (Charpinet et. al. 2001). Data on 30 of such pulsators, also known as sdB variables (sdBVs), are reported by Charpinet (2001), including the sdBV HS0702+6043, the detailed results of which are reported in Dreizler et al. (2002), and the two sdBVs PG1325+101 and PG2303+019, recently discovered by Silvotti et al. (2002). The 31st star of this class was discovered by Piccioni et al. (2000) and later confirmed by Ulla et al. (2001). These stars show short period ($1-10\\,\\rm min$) and low amplitude (1-50 milli-magnitude [mma]) non-radial pulsation modes. Charpinet et al. (1996,1997) have shown that the non-radial pulsations are probably driven by an opacity bump associated with iron ionization.  Here we report the discovery of a new variable of this class, the 32nd one, namely the PG~1613+426 sdB, thanks to the data collected at the Catania Astrophysical Observatory. This star, with B = 14.14 mag, is brighter than the average of the presently known sdB pulsators. This object has been selected from the list of ultraviolet-excess Palomar Green objects (Saffer et. al 1994), where a determination of the main atmospheric parameters is available. For this star, low resolution spectroscopy ($\\sim$~6\\,\\AA) leads to a sdOA spectral classification due to its strong, broad Balmer and HeI absorption lines, and in particular gives $T_{\\rm eff}=34\\,400\\,\\pm\\,500\\,{\\rm K}$, $\\log g = 5.97 \\pm 0.12$, and $\\rm N(He)/N(H)=0.022\\,\\pm\\,0.003$ from a least-square fitting of Balmer lines with LTE models. These characteristics place this objects near the hot end of the theoretical sdBVs instability strip. ",
        "conclusions": ""
    },
    "0212/astro-ph0212486_arXiv.txt": {
        "abstract": "We model jets in low-luminosity (FR\\,I) radio galaxies as intrinsically symmetrical, axisymmetric, decelerating relativistic flows with transverse velocity gradients. This allows us to derive velocity fields and the three-dimensional distributions of magnetic-field ordering and rest-frame emissivity. A conservation-law analysis, combining the kinematic model with X-ray observations of the surrounding IGM, gives the profiles of internal density, pressure, Mach number and entrainment rate along the jets. We summarize our recently-published results on 3C\\,31 and outline new work on other sources and adiabatic jet models. ",
        "introduction": "\\label{Intro}  This paper is a progress report on a project whose aim is to derive quantitative estimates for the physical parameters of jets. We have developed techniques to derive the three-dimensional distributions of velocity, rest-frame emissivity and magnetic-field structure and hence to deduce the jet dynamics via a conservation-law approach.  Our fundamental assumption is that jets may be approximated as intrinsically symmetrical, time-stationary, axisymmetric relativistic flows. We model their observed radio synchrotron emission in total intensity and linear polarization, using the observed differences between approaching and receding jets to constrain velocity, emissivity and field. We then combine this model of jet kinematics with a description of the surrounding IGM and use conservation of energy, momentum and particles to estimate the internal pressure and density.  For this technique to work, we need sources with two-sided but asymmetrical and straight radio jets. We then make deep observations with many resolution elements to derive total intensity and linear polarization, corrected for any Faraday rotation. We also require X-ray observations of the hot gas surrounding the radio source from which the external pressure and temperature may be derived. At present, these requirements are met only by VLA and {\\sl Chandra} observations of nearby FR\\,I radio galaxies. ",
        "conclusions": ""
    },
    "0212/astro-ph0212353_arXiv.txt": {
        "abstract": " ",
        "introduction": "Supernova remnants and their associated pulsars are sites in which strong shocks act to accelerate particles to extremely high energies. The connection between SNRs and the energetic cosmic rays that pervade the Galaxy has long been assumed, for example; shock acceleration by the SNR blast wave provides ample energy for the production of multi-TeV particles, and recent observations of nonthermal X-ray and VHE $\\gamma$-ray emission from several SNRs has confirmed the presence of electrons at these high energies. At the same time, models for the structure of PWNe predict particle acceleration at the wind termination shock, and recent X-ray observations have begun to probe these acceleration sites.  These X-ray measurements provide constraints on the expected higher energy emission processes, and are thus of considerable interest in the context of VHE $\\gamma$-ray astronomy.  Here I review recent {\\em Chandra} studies of nonthermal X-ray emission in SNRs and PWNe, with particular emphasis on the evidence for particle acceleration to very high energies. ",
        "conclusions": "The strong shocks in SNRs and PWNe have long been regarded as sites where particles are accelerated to extremely high energies. SNRs, in particular, are leading candidates as the source of cosmic rays up to the knee of the spectrum. X-ray emission from PWNe also require particles with higher energies than expected in the postshock flow of the pulsar wind, indicating efficient acceleration at the termination shock. X-ray observations have now begun to provide direct evidence of the energetic particles and shock structures where this such acceleration takes place. Through studies of the dynamics and nonthermal X-ray emission from SNRs, and of the wind termination shocks and associated particle outflows in PWNe, strong constraints are being placed on models of the particle acceleration process. High resolution X-ray observations promise continued advances in this area, through measurements of SNR nonthermal emission and expansion rates, and the inner structure of PWNe. In addition, $\\gamma$-ray observations with current and upcoming \\v{C}erenkov telescopes, as well as future space-borne observatories, hold considerable promise for probing these sites of extremely energetic particles."
    },
    "0212/astro-ph0212023_arXiv.txt": {
        "abstract": "}[2]{{\\footnotesize\\begin{center}ABSTRACT\\end{center} \\vspace{1mm}\\par#1\\par \\noindent {~S}{\\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 }} { The microlensing event OGLE-2002-BLG-055 has a single, but very reliable data point, deviating upward from a single source microlensing light curve by 0.6 mag.  The simplest interpretation calls for a binary lens with a strong parallax effect and the mass ratio in the range 0.01 - 0.001, putting the companion in the Jupiter mass range.  Given only a single deviant point it is impossible to fit a unique model.  We propose a modification of OGLE observing strategy: instant verification of a reality of future deviant points, followed by a frequent time sampling, to make a unique model fit possible. }{~}  \\noindent {\\bf Key words:}{\\it Dark matter - Gravitational lensing - Planets} ",
        "introduction": "%  Mao and Paczy\\'nski (1991) and Gould and Loeb (1992) proposed a search for planets using gravitational microlensing.  The results so far have been inconclusive (e.g. Bennett et al. 1999, Rhie et al. 2000, Albrow et al. 2000, Gaudi et al. 2002). One of us (Jaroszy\\'nski 2002) analyzed 18 candidate binary microlensing events from the catalog of OGLE-II microlensing events (Wo\\'zniak et al. 2001).  In two cases: SC20\\_1793 and SC20\\_3525 well fitting models with extreme mass ratios were found, indicating a possibility of planetary events. Unfortunately, neither case was very strong, as alternative models were also possible.  In 2001 OGLE-III (Udalski et al.  2002) begun its operation, and almost 400 candidate events were detected toward the Galactic Bulge in the 2002 observing season:  \\centerline{http://www.astrouw.edu.pl/$\\sim$ogle/ogle3/ews/ews.html} \\noindent Among them a number of events showed isolated measurements deviating a lot from otherwise smooth light curves.  While some of them were due to various instrumental (or cosmic ray) effects, some might be real. In particular, the event OGLE-2002-BLG-055, as shown in Fig. 1, had a single data point which was 0.6 mag `too bright'.  The reality of this measurements was verified on the CCD image by Dr. A. Udalski (private communication).  The fact that nearby points do not deviate appreciably suggests that this may be a binary event with an extreme mass ratio, i.e. it may be an evidence for a planet.  \\begin{figure}[h] \\center{ \\includegraphics[height=90mm,width=90mm]{fig1.eps}% } \\caption{\\small The event 2002-BLG-055 as presented in the OGLE website with the model light curve. } \\end{figure} ",
        "conclusions": "%  Given only one strongly deviating data point it is not possible to determine a unique mass ratio for a binary lens model.  It may also be difficult to persuade skeptics that a single point should be treated as a proof that the lens is binary.  Therefore, rather than argue about OGLE-2002-BLG-055 we propose a modest modification of the OGLE Early Warning System (EWS, Udalski et al. 1994).  The current system alerts on a time scale of several days, as it is tuned to stellar lenses, and it currently monitors light variations of about 150 million stars on a roughly 24 hour cycle.  At any given time there are only several hundred recognized microlensing event, which is less than 1/100,000 of all stars.  It should be relatively easy to intercept the data related to these objects and to process them within minutes of the CCD exposure and to present them to the observer in a form of a standard EWS light curve, as shown on the WWW.  In principle a data point deviating from the model fitted to the lensing event could be identified with software, but there is no urgent need for that: we have identified only 27 strongly deviant points among nearly 400 light curves of microlensing events detected in 2002, i.e. one per week.  All these points were less extreme than the one in OGLE-2002-BLG-055, and very likely most of them are not real anomalies.  Upon detecting a possible anomaly the observer may interrupt the prescheduled observing procedure and take a second exposure of the interesting field.  If the anomaly is confirmed - the observing schedule could be interrupted to provide frequent  measurements of the interesting event.  Figure 3 in this paper and our simulations indicate that extra observations should be sufficient to distinguish between models with various mass ratio. Our estimate of the possible caustic crossing time for the event OGLE-2002-BLG-055 suggests that it would be enough to make observations once per hour.  Note, that given proper sampling of a planetary event should provide a fairly easy determination of the mass ratio.  A fit to the stellar event provides the time scale $ t_E $, and the impact parameter $ u_0 $.  The timing of a short duration anomaly provides information about an approximate location of the planet in the lens plane, and the mass ratio is the only parameter that has to be thoroughly searched for.  The event OGLE-2002-BLG-055 exhibited a strong parallax effect. Had the rapid EWS system been in place there is a good chance that a caustic crossing due to planetary lensing would have been detected, and the source would have been resolved, i.e. the relative proper motion would have been determined.  Combining all three elements: the stellar event time scale $ t_E $, the parallax effect, and the relative proper motion, would allow not only to determine the mass ratio of the system, but also the masses.    \\Acknow{This work was supported with NASA grant NAG5-12212 and NSF grant AST-0204908.  We are very grateful to Dr. A. Udalski for verification of the reality of the very bright data point, and to Dr. S. Mao for providing the binary lens software.  It is a great pleasure to acknowledge that all figures were made using the SM plotting package provided by R. H. Lupton.}"
    },
    "0212/astro-ph0212509_arXiv.txt": {
        "abstract": "Using the EGRET data and an improved point source analysis, including an energy dependent point spread function and an unbinned maximum likelihood technique, we have been able to place considerably lower limits on the gamma ray flux from the galactic center region. We also test this method on known sources, the Crab and Vela pulsars.  In both cases, we find that our method improves the angular precision of EGRET data over the 3EG catalog.  This new limit on gamma rays from the galactic center can be used to test models of annihilating supersymmetric dark matter and galactic halo profiles. We find that the present EGRET data can limit many supersymmetric models if the density of the galactic dark matter halo is cuspy or spiked toward the galactic center. We also discuss the ability of GLAST to test these models. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/astro-ph0212215_arXiv.txt": {
        "abstract": "High-precision spectrophotometry is highly desirable in detecting and characterizing close-in extrasolar planets to learn about their makeup and temperature.  For such a goal, a modest-size telescope with a simple low-resolution spectroscopic instrument is potentially as good or better than a complex general purpose spectrograph since calibration and removal of systematic errors is expected to dominate. We use a transmission grating placed in front of an imaging CCD camera on Steward Observatory's Kuiper 1.5 m telescope to provide a high signal-to-noise, low dispersion visible spectrum of the star HD 209458. We attempt to detect the reflected light signal from the extra-solar planet HD 209458b by differencing the signal just before and after secondary occultation. We present a simple data reduction method and explore the limits of ground based low-dispersion spectrophotometry with a diffraction grating.  Reflected light detection levels of 0.1\\% are achievable for $5000-7000$~\\AA, too coarse for useful limits on ESPs but potentially useful for determining spectra of short-period binary systems with large $(\\Delta m_{vis}=6)$ brightness ratios. Limits on the precison are set by variations in atmospheric seeing in the low-resolution spectrum. Calibration of this effect can be carried out by measurement of atmospheric parameters from the observations themselves, which may allow the precision to be limited by the noise due to photon statistics and atmospheric scintillation effects. ",
        "introduction": "The indirect detection of large hot Jupiter-like extrasolar planets with periods on the order of days raises the possibility of direct detection of the reflected light from the planet's atmosphere. The expected contrast ratio for these close-in gas giant planets is expected to be quite favorable $(10^4-10^5)$ compared to the contrast ratio between Jupiter and our sun $(10^{9})$.  However, the close proximity of these planets to their stars precludes spatially resolving the planet with current telescopes. Thus additional information about these planets requires other methods of discerning its flux in the presence of the much brighter star. Several approaches are possible to separate the starlight from the planet's light.  \\cite{cha99} and \\cite{col01} each observed tau Boo with a high resolution spectrometer in an attempt to discern the planet's light through by the Doppler shift of the reflected light as the planet orbits its star.  Although no detection was observed, the data were sufficient to place upper limits on the planet's albedo.  The discovery that the planet around HD 209458 occulted its star during the 3.5 day orbit \\citep{cha00} was a significant step in verifying the existence of the extrasolar planet candidates. \\cite{bro01} followed up this initial detection with high precision HST photometric measurements of the transit. This yielded a more precise light curve of the occultation, and gave the first indication of the composition of the planet's atmosphere. The transit was detected to be slightly deeper at the wavelength of sodium absorption \\citep{cha02}, as predicted by various groups \\citep{bro01a,sea00,hub01}.  Direct detection of the reflected light from a close-in planet has the potential of telling us the albedo of a planet, and thus also, the temperature of the planet.  From a more pedestrian viewpoint detection of the planet's reflected light would be the indication of what the planet ``looks'' like. The color is more than just a curiosity, though, allowing the first constraints on the planet's chemical makeup and vertical structure (such as the existence of clouds).  We observe the HD209458 system just before and after the eclipse of the planet by its parent star.  Theoretical models suggest the planet may have a bright (0.7) albedo with a wide ($R/\\Delta R\\sim 20$) absorption feature due to pressure broadened sodium at 591 nm \\citep{sud00}.  The reflected light fraction from the star is $2\\times10^{-4}$ for a geometric albedo of one, so a detection level of $\\simeq 1\\times10^{-4}$ per wavelength band is needed to produce a useful constraint on the albedo of the planet.  Colour information is obtained with the use of a transmission grating in front of a CCD detector, in an objective prism configuration. Objective prism surveys are used with wide field imaging telescopes for spectral typing of stars and classifications of the brighter galaxies. The low dispersion, high throughput and no slit jaw losses make the objective prism method suitable for colour dependent transits. The main challenge is then in obtaining accurate spectrophotometry with the variation in transmission of the atmosphere as a function of airmass and the variation of spectral resolution due to the local seeing at the telescope. ",
        "conclusions": "The current approach to obtaining data of HD~209458 results in very inefficient observing. For every 5 second exposure, approximately 30 seconds of real time is needed in order to read out the CCD, resulting in 16\\% efficiency to the observing. The exposures are limited to 5 seconds currently to avoid saturating the CCD.  However, an increase in the efficiency can be made by spreading the light out perpendicular to the direction of dispersion with a cylindrical lens.  This will allow much longer exposures resulting in an efficiency $>50\\%$. This would result in a decrease in the noise associated with both the photon and scintillation noise allowing reduction of these quantitie to well below the expected signal due to a planetary secondary eclipse.  The power of the technique we describe is that it requires very little in extra calibration data - an estimate of the atmospheric blurring is made from transverse cuts across the spectrum, and each frame thus provides its own calibration data.  Precision spectrophotometry from ground-based telescopes is an attractive technique for probing secondary eclipses of planetary orbits, but the current work shows the technique is limited by systematic effects of atmospheric extinction to precision of 0.1\\%. It is possible that future refinements of the technique will allow an improvement by calibrating out the effect of the atmosphere through appropriate standard star observations.  This will allow explorations of intensity variations down to the level set by the photon and scintillation noise of the star."
    },
    "0212/astro-ph0212201_arXiv.txt": {
        "abstract": "{ Hipparcos observations of some variable stars, and especially of long-period (e.g. Mira) variables, reveal a motion of the photocenter correlated with the brightness variation ({\\it variability-induced mover} -- VIM), suggesting the presence of a binary companion.  A re-analysis of the Hipparcos photometric and astrometric data does not confirm the VIM solution for 62 among the 288 VIM objects (21\\%) in the Hipparcos catalogue.  Most of these 288 VIMs are long-period (e.g. Mira) variables (LPV). The effect of a revised chromaticity correction, which accounts for the color variations along the light cycle, was then investigated. It is based on `instantaneous' $V-I$ color indices derived from Hipparcos and Tycho-2 epoch photometry. Among the 188 LPVs flagged as VIM in the Hipparcos catalogue, 89 (47\\%) are not confirmed as VIM after this improved chromaticity correction is applied. This dramatic decrease in the number of VIM solutions is not surprising, since the chromaticity correction applied by the Hipparcos reduction consortia was based on a fixed $V-I$ color. Astrophysical considerations lead us to adopt a more stringent criterion for accepting a VIM solution (first-kind risk of 0.27\\% instead of 10\\% as in the Hipparcos catalogue). With this more severe criterion, only 27 LPV stars remain VIM, thus rejecting 161 of the 188 (86\\%) of the LPVs defined as VIMs in the Hipparcos catalogue. ",
        "introduction": "The Hipparcos satellite \\citep{Hipparcos} obtained high-precision photometric and astrometric data for some 120\\,000 stars, discovering in the process many new binary systems.  In addition, a small number of possible binaries (288) were identified as variability-induced movers \\citep[hereafter VIM,][]{Wielen-1996}.  The physical origin of this effect lies in a close binary companion of an intrinsic variable star, causing the photocenter of the system to move as the primary's luminosity varies.  The Hipparcos analysis allowed for the presence of such a companion, solving for its position with respect to the primary in addition to the position, parallax and proper motion of the system.  Although the VIM model is likely to be appropriate in some cases \\citep{Bertout-1999:a}, the results for some other systems are questionable.  For instance, rather few VIM have been resolved by speckle interferometry even though they have been extensively observed by, for instance, the Washington team \\citep{Mason-1999:b,Mason-2001:a}.  This situation led us to re-investigate the VIM solutions using an improved processing of the Hipparcos data.  The entire available set of astrometric data was used, instead of only FAST data, as was the case with the preparation of the Hipparcos catalogue. This results in a decrease by 21\\% of the number of VIM solutions, especially the red ones (Sect.~\\ref{Sect:OrigVIM}).  During processing for the Hipparcos catalogue,  corrections were applied to the Hipparcos photometry and astrometry to correct for the intrinsic chromaticity of the optical system (i.e. the position of the diffraction spot on the detector depends upon the color of the star) and for the aging effects of the  optics and detectors during the mission. These effects are especially important for red variables because of their extreme red colors and because their variability time scales are comparable to the duration of the mission. Moreover, red variables have changing colors, and this variation must be taken into consideration when applying the chromaticity correction. In the preparation of the Hipparcos catalogue, the reduction consortia nevertheless adopted a {\\em constant} color in the chromaticity correction process. There is thus room for improvement using chromaticity corrections based on {\\em epoch} colors, rather than on average colors. \\citet{Platais-2002:a} have devised a scheme to derive epoch colors based on Tycho-2 and Hipparcos epoch photometry. An improved chromaticity correction has then been applied to a sample of 188 long-period variable stars flagged as VIM in the Hipparcos catalogue (Sect.~\\ref{Sect:NewVIM}), using the epoch $V-I$ index derived from $Hp - V_{T2}$. Sect.~\\ref{Sect:Results} shows that a more stringent criterion (first-kind risk of only 0.27\\%) for accepting VIM solutions seems appropriate based on astrophysical considerations. ",
        "conclusions": "\\label{Sect:Conclusion}  Instantaneous $V-I$ indices derived from a color transformation based on $Hp - V_{T2}$ have been used to improve the chromaticity correction of red and variable stars. All stars flagged as Variability-Induced Mover (VIM) in the Hipparcos catalogue have been reprocessed. It turns out that about 85\\% of those objects can now be satisfactory modeled with the basic 5-parameter, single-star solution. This result explains why many of the Hipparcos VIM solutions did not receive confirmation of their binary nature from speckle observations."
    },
    "0212/astro-ph0212037_arXiv.txt": {
        "abstract": " ",
        "introduction": "Super Soft Sources (SSS) are characterized by radiation with effective temperatures of $\\sim$$10^{5}$ to $\\sim$$10^{6}$~K and luminosities of $\\sim$$10^{36}$ to $\\sim$$10^{38}\\ {\\rm erg}\\ {\\rm s^{-1}}$. The first luminous SSS in the Large Magellanic Cloud (LMC) have been discovered around 1980 with the {\\sl Einstein} observatory (Long et al. 1981). Optical identifications for both CAL~83 and CAL~87 established the binary nature of these sources (Cowley et al. 1990; Smale et al. 1988). During  the {\\sl ROSAT} all-sky survey in 1990--1991 in soft X-rays (cf. Tr\\\"umper et al. 1991) and pointed follow-up observations many SSS have been detected. {\\sl BeppoSAX} discovered super-soft X-ray emission from recurrent and classical novae. Deep observations performed with {\\sl Chandra} and {\\sl XMM-Newton} led to the discovery of SSS in nearby spiral and elliptical galaxies. A catalog of SSS is given in Greiner (2000a). Optical identifications exist for SSS in the Galaxy, the LMC, and the SMC (cf. Tab.~11.1).   \\begin{figure*} \\vbox{\\hskip -\\leftskip \\psfig{figure=chp11f1.ps,width=14.0cm,angle=-180.0,% bbllx=3.0cm,bblly=6.0cm,bburx=17.0cm,bbury=27.5cm,clip=}}\\par \\end{figure*}   Here binary SSS are considered. For 7 systems orbital periods of $\\sim$$9^{\\rm h}$ to $\\sim$$4^{\\rm d}$ have been determined, in the range predicted for steadily nuclear burning white dwarfs (WDs) in close-binary systems undergoing unstable mass-transfer on a thermal timescale (close-binary SSS, CBSS, DiStefano and Nelson 1996; van den Heuvel et al. 1992). From binary evolutionary calculations follows that such systems may have evolved companions with masses $\\sim$$(1.3-3.0)\\ M_{\\odot}$ (Rappaport et al. 1994). U~Sco is a recurrent nova. A few systems with orbital periods of $1.4^{\\rm h}$ to $3.5^{\\rm h}$ are classical novae (CV-type SSS): V~1974 Cyg (Krautter et al. 1996), GQ~Mus (\\\"Ogelman et al. 1993), Nova LMC 1995 (Orio and Greiner 1999), and V~382~Vel (Orio et al. 2002). 1E~0035.4-7230 and RX~J0537.7$-$7034 may have a lower mass WD. The recurrent novae RS~Oph, T~CrB, V394~CrA, and CI Aql 2000 have a subgiant or low-mass giant donor and modeling of the optical outburst light curve requires a super-soft phase (Hachisu and Kato 2000a,b; 2001a,b). The symbiotic systems with detected super-soft X-ray emission RR~Tel (Jordan et al. 1994), AG~Dra (Greiner et al. 1997), SMC~3 (Jordan et al. 1996), and Lin~358 (M\\\"urset et al. 1997) have orbital periods of at least a few hundred days (cf. Mikolajewska and Kenyon 1992). There are two transient SSS which have not yet been optically identified, the globular cluster X-ray source 1E~1339.8+2837 (Dotani et al. 1999) and RX~J0550.0$-$7151 (Reinsch et al. 1999). For a review of SSS see Kahabka and van den Heuvel (1997). ",
        "conclusions": ""
    },
    "0212/astro-ph0212347_arXiv.txt": {
        "abstract": "The status and capabilities of the KASCADE-Grande extensive air shower experiment are presented.  The installation is located at Forschungszentrum Karlsruhe and comprises a large collecting area (0.5 km$^2$) electromagnetic array (Grande) operated jointly with the existing KASCADE detectors.  KASCADE-Grande will cover the primary energy range $10^{16}$ eV $< E_0 < 10^{18}$ eV overlapping with KASCADE around $10^{16}$ eV, thus providing continuous information on the primary energy and mass of cosmic rays from $3 \\cdot 10^{14}$ eV up to $10^{18}$ eV. The major goal of the measurements is the unambiguous observation of the``iron knee'' expected in the cosmic ray spectrum at $E_k^{Fe} \\approx 10^{17}$ eV. ",
        "introduction": "The origin and acceleration mechanism of ultra-high energy cosmic rays ($E \\ga 10^{14}$ eV) and the origin of the `knee' at approx.\\ $3\\cdot10^{15}$ eV have been subject to debate for several decades.  Most commonly, it is assumed that cosmic rays gain their energy by first order Fermi acceleration at highly supersonic shock waves originating in supernova explosions. Simple dimensional estimates show that this process is limited to $E_{\\rm max} \\la Z \\times (\\rho \\times B)$, with $Z$ being the atomic number of the cosmic ray (CR) isotope and $\\rho$, $B$ the size and magnetic field strength of the acceleration region.  A more detailed examination of the astrophysical parameters suggests an upper limit of acceleration of $E_{\\rm max} \\approx Z\\times 10^{15}$\\,eV \\cite{drury94b,berezhko99}.  It is thus tempting to identify the knee with the maximum energy obtained by such accelerators.  A brief compilation of alternative interpretations of the knee is given in Ref.~\\cite{kampert01}.  As a simple consequence of the aforementioned `standard' picture, one expects to see a break at constant rigidity for all particles, i.e.\\ protons would break off first and iron nuclei at correspondingly higher energy.  This results in an increasingly heavier composition at energies above the knee.  Such a global change of composition is in fact observed by various experiments, including KASCADE \\cite{glasstetter99,antoni02} and EAS-TOP \\cite{aglietta99} (a recent compilation can be found in Ref.~\\cite{swordy02}).  A rather convincing proof has recently been presented by the KASCADE collaboration \\cite{ulrich01,kampert01} reporting preliminary evidence for a knee shifting towards higher energies with increasing mass of the primaries (see also contributions to this conference by H. Ulrich {\\it et al.} and M. Roth {\\it et al.}).  However, reconstruction properties and statistics do not allow to clearly identify the expected iron knee at $E \\approx 10^{17}$ eV. The primary goal of the KASCADE-Grande experiment is to uncover the possible existence of an iron break by providing high quality data from $10^{16}$ - $10^{18}$ eV. The experiment covers an area of approx.\\ 0.5~km$^2$ and is located next to the KASCADE site in order to operate jointly with the existing KASCADE detectors.  The combination with KASCADE will guarantee a cross calibration of the detectors and will provide full coverage from $3 \\cdot 10^{14}$ to $10^{18}$ eV.  By the time of writing this article, all detector stations have been deployed and first showers were already observed with preliminary readout electronics.  Full operation will start in January 2003 and full statistics will be reached after 3-4 years of data taking.  The present status of the array, its resolutions and capabilities in reconstructing the average primary composition and its verification of the hadronic interaction models are discussed. ",
        "conclusions": "KASCADE-Grande has been realized at the Forschungszentrum Karlsruhe by a joint operation of the KASCADE and EAS-TOP detectors.  The apparatus is almost completed and data taking will start shortly.  Within three years of effective data taking it will accumulate sufficient statistics up to $10^{18}$ eV. The experiment will cover an energy range that is only poorly known, essentially from old AKENO \\cite{as84} and the lower tail of FLY'S EYE \\cite{as94} data.  Its task is to provide more information on the structure of the knee by testing the rigidity dependent model up to the energies of the ``iron'' group. Moreover, it will allow to test hadronic models in an energy range not accessible to present accelerators but important for CRs physics.  Finally, it will provide a bridge to the measurement and interpretation of CRs for experiments working at much higher energies like the Pierre Auger, HIRES or EUSO projects.  \\subsection*{Acknowledgements} The authors are indebted to the members of the engineering and technical staff of the KASCADE-Grande collaboration, who contributed with enthusiasm and engagement to the construction of the experiment.  The work has been supported by the Ministry for Research and Education of the Federal Government of Germany. Special thanks also go to INFN for allowing the use of the EAS-TOP equipment to build the Grande array.  The collaborating Institute of Lodz is supported by the Polish State Committee for Scientific Research and the Institute of Bucharest by a grant of the Romanian National Academy for Science, Research and Technology.  The KASCADE-Grande work is embedded in the frame of scientific -technical cooperation (WTZ) projects between Germany and Romania, and between Germany and Poland."
    },
    "0212/astro-ph0212171_arXiv.txt": {
        "abstract": "{ Non-LTE abundances of magnesium, aluminum and sulfur are derived for a sample of 23 low-$v \\sin i$ stars belonging to six northern OB  associations of the Galactic disk within 1 kpc of the Sun. The abundances are obtained from the fitting of synthetic line profiles to high resolution spectra. A comparison of our results with HII region abundances indicates good  agreement for sulfur while the cepheid abundances are higher. The derived abundances of Mg show good overlap with the cepheid results. The aluminum abundances for OB stars are significantly below the cepheid values. But, the OB star results show a dependence with effective temperature and need further investigation. The high Al abundances in the cepheids could be the result of mixing.  A discussion of the oxygen abundance in objects near the solar circle suggests that the current mean galactic oxygen abundance in this region is 8.6-8.7 and in agreement with the recently revised oxygen abundance in the solar photosphere. Meaningful comparisons of  the absolute S, Al and Mg abundances in OB stars with the Sun must await a reinvestigation of these elements with 3D hydrodynamical model atmospheres for the Sun. No abundance gradients are found within the limited range in galactocentric distances in the present study. Such variations would be expected only if there were large metallicity gradients in the disk. ",
        "introduction": "The massive OB stars, as well as other young objects such as cepheids or H~II regions, are used commonly as tracers of the current chemical composition in the Galactic disk.  Abundance analyses of H~II regions are restricted to a handful of elements, those being most typically He, N, O, S, Ne, and Ar.  In the OB stars, besides N, O, and S, the additional elements C, Mg, Al, Si, and Fe can be analyzed using both LTE and non-LTE techniques. Cepheids are the  more evolved, cooler descendents of a subset of the OB stars  (with masses of $\\sim$ 5-10M$_{\\odot}$) and a rather large number  of elements ($\\sim$ 25) can be detected in their spectra.  As all three of these types of objects are young ($\\le$ few $\\times$ 10$^{7}$ yr), it is reasonable to expect that examples of them inhabiting the same region of the Galaxy should contain approximately the same mixture of chemical elements.  This approximation should be as good as allowed by small scale chemical inhomogeneities that can be produced on timescales of the lifetimes of large, starforming regions, i.e. a few times 10$^{7}$ yr.  Care must be taken in comparing abundances between H~II regions, OB stars, and cepheids, however. For example, internal stellar mixing may affect both OB stars and cepheids to varying degrees. Some evidence of the mixing of material exposed to the CN-cycle has been uncovered in certain OB stars, even near the main sequence (Gies \\& Lambert \\cite{gl92}); this mixing may be driven primarily by rotation (Heger \\& Langer \\cite{hel00}).  In this case, nitrogen abundances are measurably enhanced (by $\\sim$ +0.3 dex)  with, perhaps, a marginal decrease in the carbon abundance (by $\\sim$ $-$0.1 dex), while oxygen remains untouched.  In the more evolved cepheids, even deeper mixing may have occurred that involves the full CNO cycles, such that some oxygen depletion might be detectable (as well as larger N enhancements and C depletions). There is also the possibility that surface abundances of sodium, magnesium, and aluminum have been altered by the mixing of material exposed to the Ne-Na and Mg-Al cycles.  In the H~II regions, uncertainties or systematics may arise from temperature fluctuations in the gas, unknown radiation environments, or possible depletions of some elements out of the gas phase and onto solid particles.  By concentrating on a set of OB stars, in comparison to sample H~II regions and cepheids, all near the solar circle, abundance trends found in these various objects can be intercompared and ultimately be checked against their corresponding solar values. Such checks can help uncover possible inconsistencies and provide stronger constraints on abundance gradients derived from different types of objects.  That is the aim of this paper, which is the fourth paper in a series whose ultimate goal is to use a large sample of OB stars to trace abundance gradients in the Galactic disk.  Paper I (Daflon, Cunha \\& Becker \\cite{pap1}) presented the first results from this survey  and concentrated on 8 sharp lined star ($v \\sin i\\le$ 60 km s$^{-1}$) members of the Cep~OB2 association: LTE abundances of C, N, O, Si, and Fe were derived, as well as non-LTE abundances of C, N, O, and Si.  Daflon et al. (\\cite{pap2} $-$ Paper II) added analyses of 15 members of five additional northern OB associations (Cep~OB3, Cyg~OB2, Cyg~OB7, Lac~OB1, and Vul~OB1).  The same sets of elements were analyzed in non-LTE as in Paper I, but additional LTE results were presented for Mg, Al, and S.  In Paper III  (Daflon et al. \\cite{pap3}), the analysis was expanded to include some of the more rapidly rotating stars ($v \\sin i$= 60-150 km s$^{-1}$) in the northern sample.  In addition, the atomic analysis included non-LTE calculations for Mg and Al. In this paper, the non-LTE calculations for Mg and Al are applied to the remaining northern stars, and sulfur is now added as an element that can be studied in non-LTE. ",
        "conclusions": "We have presented non-LTE abundances of magnesium, aluminum and sulfur for 23 OB stars members of OB associations within 1 Kpc of the Sun. Magnesium abundances derived for cepheids in the literature agree well with the results for OB stars while no Mg abundances can be derived for H II regions.  On the other hand, sulfur results point to discrepencies between the cepheids on one side and the OB stars and H~II regions on the other, that are not easily resolved: the cepheid sulfur abundances are derived from S~I lines, while the OB stars use S~III lines, and the H~II region results come from far-IR [S~III] lines at 19 $\\mu$m.  Aluminum abundances OB stars  must be interpreted with  caution: the relatively low Al abundances derived may reflect problems in the non-LTE calculations for this species and deserves additional analysis before further conclusions while Al in the cepheids could be enhanced due to internal mixing."
    },
    "0212/astro-ph0212421_arXiv.txt": {
        "abstract": "We present a semianalytic model of steady-state magnetically and radiatively driven disk outflows in Active Galactic Nuclei (AGNs) consisting of a continuous wind with embedded clouds.  The continuous outflow is launched from the disk surface as a centrifugally driven wind, whereas the clouds are uplifted from the disk by the ram pressure of the continuous outflow.  The inner regions of the outflow shield the outer portions from the strong ionizing central continuum, enabling gas in the outer regions to be radiatively accelerated.  In this paper, we describe the model in detail, outline the tests used to verify its accuracy, and then compare it to other outflow models already in use, showing that some previous AGN wind simulations do not consider radiative transfer in the detail necessary for accurate modeling. We also perform a comprehensive parameter study to explore the dependence of the continuous wind properties on the relevant physical parameters.  We find that the introduction of clouds has a significant impact on the geometry, density, and terminal velocities of the continuous wind: in particular, changes of up to 50\\% in the wind velocity can be directly attributed to the effect of clouds. Within this survey, line driving is often the dominant acceleration mechanism for clouds, owing to their small volume filling factor, which yields an overall low cloud opacity.  Continuum driving can also become important for dusty clouds.  For the continuous wind alone, we find that radiative acceleration can have a significant impact on both the terminal velocities and the kinematics near the surface of the disk. However, unlike the clouds, the radiative acceleration of the continuous wind comes largely from continuum opacities, not from line driving as in some previous models.  Line driving produces comparable acceleration only for parts of the wind far from the central source, where the ionization state of the gas drops due to the lower incident flux.  When the winds are dusty, their terminal velocities change by approximately 20\\% to 30\\% (depending on the type of dust) over pure centrifugally-driven outflows. ",
        "introduction": "Active Galactic Nuclei (AGNs) often declare themselves boisterously to observers: well-known for their incredibly bright nuclei, they can also display emission line widths of up to several thousand ${\\rm km~s}^{-1}$, and absorption lines in the optical, UV, and X-ray, with blueshifts up to 60,000 ${\\rm km~s}^{-1}$.  They do not yet, however, communicate as well with theorists: we lack a physical picture of the geometry and kinematics of gas in the core of AGNs.  Researchers agree that AGNs house supermassive black holes, but other components of the theoretical picture spark heated debate.  The payoff for building a successful model of the geometry and kinematics of gas in AGNs is very high: not only will we gain understanding of how AGNs work, and insight into the almost ubiquitous phenomenon of accretion and the physics of black holes, but we may also learn about how galaxies evolve \\citep[e.g.,][]{Blandford2001}.  In the absence of any well agreed-upon theoretical model, researchers have put together a phenomenological picture of the structure of AGN, called the ``Unification Model''\\citep{Ant93, UP95}.  As commonly cited, this paradigm consists of a supermassive black hole with an accretion disk, as well as a coplanar ``torus'' of gas and dust that orbits near the outskirts of the accretion disk; the location, composition, and dimensions of the torus remain unknown.  Gas, commonly modeled as clouds and postulated to exist above the accretion disk, produces both the broad emission lines (within the Broad Emission Line Region, or BELR) and, further from the central source, the narrow emission lines.  This model has proven itself useful through its explanation of the difference between AGNs that contain only narrow emission lines, and those that contain both broad and narrow lines: if the AGN is oriented so that our line of sight intersects the torus, it obscures the broad lines region, and we see only narrow emission lines.  In recent years, models that further elucidate the geometry and kinematics of gas in the unification model have also been proposed \\citep[e.g.,][]{Elvis00,Ganguly01}; although these models cite the need for such mechanisms as radiative acceleration, they remain chiefly empirical.  Observational hints are now appearing that the unified, ``inclination angle'' paradigm for AGNs may not work in all objects.  In fact, relatively few AGNs have been used as proof of this unification model \\citep{Axon01}.  In addition, inclination-angle independent measures of various AGNs seem to show variations which are not in agreement with the unified model \\citep{Weymann02}.  Recent X-ray observations of AGNs seem to show objects that have large X-ray absorbing columns, and yet have broad optical emission lines, while other sources with only narrow optical emission lines have no apparent X-ray absorption \\citep{Pappa01}.  Thus, the empirical models of gas distribution in the central regions of AGNs might not be as widely applicable as previously thought.  At the root of this difficulty lies the fact that we are still quite ignorant of what forces dominate in different parts of the nuclei of active galaxies, or what forces may dominate in different AGNs.  To solve this problem, and develop consistent dynamical models of AGNs, researchers have appealed to two major sources of acceleration: magnetohydrodynamic driving of clouds \\citep[e.g.,][]{Emmer92,Bot97}, and radiative acceleration of continuous winds \\citep[e.g.,][]{MCGV95,Pro00}.  One can also see in the above very short summary that the field is somewhat divided on whether we are watching continuous outflows or ``clumpy'' outflows [however, the continuous wind simulations of \\citet{Pro00} can produce condensations that may be the source of a multiphase medium].  Both of these types of models have had successes in modeling AGN observations; for instance, both the magnetohydrodynamic model \\citep{Bot97} and the radiative acceleration model \\citep{MC97} have reproduced the observed single-peak broad emission lines observed in AGNs.  Since these models have fit the basic observational traits of AGNs, there seems to be no clear way to select which physics may be dominant in which AGNs, and we are left with two very different outflow scenarios.  The challenge remains, therefore, to develop tools that can distinguish between these models and that will allow us to test for the presence of individual forces or components.  A general model seems necessary: one that self-consistently handles all of the elements that have been previously proposed, so that we can systematically test one inclusive model against the observations.  Such a model may help us gain insight into how AGN winds work, how gas is distributed in the cores of AGNs, and may aid in testing theories about accretion disks and AGN evolution.  This paper presents such a tool, including all of the above mentioned elements and their interactions (such as the interplay between magnetic forces and radiative driving, and the drag forces between the continuous wind and the clouds) to provide a physically self-consistent ``platform'' from which to test all of the different components, together.    This paper outlines our motivation for such a model, defines the particular construction we have chosen, and introduces a parameter survey to aid intuition of how the included processes interact.  The paper begins with an outline of past models in \\S\\ref{modelOverview}, ending with an overview of our model.  In \\S\\ref{detailedModel}, we describe the various components of our model, and in \\S\\ref{Results} we present early results from our calculations, showing the dependences of wind velocities and densities on various input parameters.  We summarize our findings in \\S\\ref{Conclusions}. ",
        "conclusions": "We have presented a self-consistent, semianalytic, steady-state disk-wind model that combines magnetic and radiative accelerations and includes two gas phases: a magnetized continuous wind and embedded diamagnetic clouds.  The continuous wind is driven centrifugally and, after rising above the disk surface, becomes subject to radiative acceleration by the central continuum source.  The clouds are uplifted from the disk surface by the ram pressure of the continuous wind, which also confines them by its (largely magnetic) pressure, and they are subsequently pushed by the radiation pressure force of the central continuum.  We calculate the radiative acceleration on both of those components from first principles, using detailed photoionization simulations to determine the ionization state of the gas and the continuum incident on that gas.  Since the model includes magnetic launching, the outflow automatically incorporates a shielding column of gas that can attenuate the central continuum and allow efficient radiative acceleration in the outer parts of the wind.  This addresses the problem of how to launch a shielding column in areas of the wind where radiative acceleration is not effective, because of the high degree of ionization of the gas.  We have used this model to illustrate the dependence of a two-phase magnetically and radiatively driven wind on various outflow parameters, and to demonstrate the care with which one must treat radiative acceleration calculations.  We find that {\\it the wind structure is significantly altered} when we include clouds that are also radiatively accelerated: their low filling factor, low ionization parameters, and the possibility of dust in the clouds all allow radiative acceleration when the continuous wind is too highly ionized for radiation to be an effective driving source.  Not unexpectedly, clouds have the largest impact when their mass outflow rate is comparable to that of the wind itself.  We have also determined that including dust in the continuous outflow can have a large impact on wind geometry and kinematics because of the increase in continuum opacity.  Dust has a large enough impact on wind structure that even changing the type of dust results in significant changes in the outflow.  Finally, we have found that higher-density winds with low shielding can be influenced heavily by radiative acceleration.  In this case, radiative acceleration is not acting through the clouds but instead comes from the continuum driving in the continuous wind itself; large changes in the wind occur with dust present in the wind, but even the gas continuum opacity can result in large radiative acceleration of the continuous wind.  This is a recurring theme in our results: continuum driving by and large plays a more important role in continuous wind acceleration than line driving.  Line driving, on the other hand, is more important in the clouds.  We have also found that, in the context of our particular cloud model, much of the spectral dependence of radiative acceleration is largely erased, with different central continua producing very similar perturbations to the wind's structure and dynamics.  Furthermore, we showed that variations in the mass of the central black hole and in the launching radius of the wind do not greatly modify the initial centrifugally-driven wind's structure or kinematics, although it is important to remember that changing both of these parameters affects the Keplerian velocity at the base of the wind, and hence changes the magnitude of the wind velocity.  Finally, we established that the velocity structure near the surface of the disk is very sensitive to radiative acceleration, which may be an important key to observationally distinguishing between pure MHD-driven winds and radiatively driven winds.  While both accelerate rapidly from the disk, radiation pressure causes much greater initial acceleration as well as much larger drops in the density as the wind lifts off from the surface of the disk.  This effect would be much more noticeable in AGNs with low shielding columns, where radiative acceleration is much stronger near the surface of the disk.  As a result of mass conservation, the variations in density with altitude above the disk closely mirror those of the velocity, and may also be valuable tracers to discriminate between wind acceleration processes. This may hint that if Broad Emission Lines are formed near the disk surface, observations of those lines may be capable of distinguishing between different acceleration processes.  It is important to note that we have not yet included the radiative acceleration due to intrinsic or reprocessed flux from the accretion disk itself as in MCGV95 and \\citet{Pro00}.  Including this effect may yield even stronger acceleration from the surface of the disk.  Our next step will be to fit this model to observations in an attempt to directly quantify the processes that act within AGN outflows.  To complete this, it will be important to widen our parameter survey to include, for example, other incident spectra and dust compositions \\citep[to examine, for instance, the dust distribution found in][]{CK01}.  We are currently developing a Monte Carlo module that will allow both line emission and absorption predictions for a range of inclination angles.  At the same time, we also hope to include the effect of radiation acceleration due to accretion disk flux and pursue observational signatures of embedded condensations that are not magnetically confined by the continuous wind \\citep[e.g.,][]{Bot01}.  In addition, we can learn a great deal about the geometry of AGN outflows from polarization measurements \\citep[e.g.,][]{Corbett00, Smith02}.  The calculation of both line and continuum polarization will also be included in our Monte Carlo program.  Finally, we note that this program is in no way restricted to considering AGN outflows alone; one can easily change the model's input parameters to study systems such as high-luminosity young stellar objects and cataclysmic variables, where similar outflows are observed."
    },
    "0212/astro-ph0212084_arXiv.txt": {
        "abstract": "I present a model-independent spherically symmetric density estimator to be used in the cross-correlation of imaging catalogs with objects of known redshift.  The estimator is a simple modification of the usual projected density estimator, with weightings that produce a spherical aperture rather than a cylindrical one. ",
        "introduction": "Measuring galaxy properties as a function of their local environment is a central task of observational extragalactic astronomy. Doing so requires a measure of density.  With redshift surveys, one can estimate densities by counting other spectroscopic galaxies in a redshift-space window around each primary object. However, imaging catalogs are generally much deeper than spectroscopic catalogs, which suggests that cross-correlating the imaging catalog with the spectroscopic catalog should yield useful densities even for the faintest spectroscopic objects.  The presence of at least one spectroscopic redshift in each pair is of great utility because it means that one can immediately map angular separations to physical distances and derive intrinsic properties (i.e.~luminosities) for both the spectroscopic object and the correlated objects from the imaging catalog.  This opportunity has been used extensively in the early study of galaxy clustering \\citep{Dav78,Yee87,Lil88,Sau92} and in the study of dwarf galaxies in groups \\citep[e.g.,][]{Fer91} and in the field \\citep{Phi87,Vad91,Lor94,Lov97}. Of course, the correlated objects are still seen in projection, with a mix of spatial separations at each transverse separation.  Angular correlations can be inverted to spatial correlations, and hence to a form of density, by the assumption of isotropy, and many of the above studies did this by assuming that the correlation functions are power laws in scale \\citep[e.g.,][]{Phi85}. \\citet{Sau92} and \\citet{Lov97} go one step further by using the assumed power-law as an optimal filter.  \\citet{Fal77} suggest that the general inversion can be done with a smoothing filter. \\citet{Bau93} use a regularized Lucy's iteration to do a similar inversion, while \\citet{Dod00} and \\citet{Eis01} use other smoothing priors.  In the study of galaxy environmental dependences on small scales, it is useful to pursue a more model-independent measure of the density. On scales of 1 Mpc, it is quite possible that the correlation functions will deviate from power laws (and certainly from a uniform power law) in ways that depend on luminosity, star-formation rate, or other variables \\citep[e.g.,][]{Scr02}.  Here I describe a method to recover density estimates in spherically symmetric real-space windows from the cross-correlation of imaging and spectroscopic catalogs independently of the shape of the correlation function.  The resulting formulae are a simple alteration of the usual background subtraction methods.  The method can be subdivided so as to yield noisy density estimates for individual spectroscopic objects that can then be averaged according to whatever subclasses one might desire. ",
        "conclusions": "I have presented a model-independent estimator for the spherically averaged overdensity of imaging catalog objects around spectroscopic objects.  The method is simple to apply; one can view it as an alteration of standard background subtraction methods to yield spherical apertures rather than cylindrical ones.  The method does not require an explicit inversion of the angular cross-correlation function to a spatial correlation function, although clearly if one finds oneself measuring the densities in multiple apertures of different sizes, one is effectively reverting back to an inversion method. The primary advantages of the new method compared to integrating the output of an inversion (i.e., computing eq.~[\\ref{eq:abel}] and then eq.~[\\ref{eq:Deltadef}]) are that one does not need a smoothing prior and that the statistic can be applied to each spectroscopic galaxy independently (eq.~[\\ref{eq:Deltaj}]), leaving one free to sum the results {\\it post facto} across as many spectroscopic subsamples as one desires. Like other angular methods, the new density estimator is unaffected by redshift distortions, which gives it an advantage on small scales over density estimation from spectroscopic catalogs.  Past work has assumed power-law correlation functions to recover spatial densities.  This allows one to achieve higher signal-to-noise ratio and hence is a better choice for some applications.  However, when probing the dependence of galaxy properties on small-scale environment, the model independence of the density estimator presented here is a valuable advantage.  With today's large surveys \\citep[e.g.,][]{Yor00,2dF}, statistical precision is sometimes less precious than systematic control.  While it is clear that one can profitably consider the dependence of average density on the properties of the spectroscopic objects \\citep{Hog02}, it is worth pointing out that one can also derive densities for subdivisions of the imaging catalog.  For example, one can probe the color and/or luminosity distribution of objects within a spherical aperture of a particular set of spectroscopic objects (e.g., galaxies or clusters). A speculative application would be to couple this approach to galaxy-galaxy weak lensing mass estimates.  If one had a mass estimate for each imaging object, assuming the spectroscopic redshift, then one could find the masses of objects that are correlated with the spectroscopic tracer."
    },
    "0212/astro-ph0212477_arXiv.txt": {
        "abstract": "{ A catalogue of extra-galactic jets is very useful both in observational and theoretical studies of active galaxies. With the use of new powerful radio instruments, the detailed structures of very compact or weak radio sources are investigated observationally and many new radio jets are detected. In this paper, we give a list of 661 radio sources with detected radio jets known to us prior to the end of December 2000. All references are collected for the observations of jets in radio, IR, optical, UV and X-ray wave-bands. ",
        "introduction": "Theoretically every extra-galactic radio source contains at its center an accretion disc surrounding a super-massive black hole and ejecting two symmetric relativistic motion plasma jets in opposite directions along its rotation axes (Begelman et al. \\cite{Begelman}). Whether a jet can be detected depends not only on the properties of the radio source itself but also on the instruments used. Radio jets have been detected in various types of radio sources: FR I and FR II radio galaxies, Seyfert galaxies, radio quasars, and BL Lac objects (Bridle \\& Perley \\cite{Bridle3}; Liu \\& Xie \\cite{Liu} hereafter Paper I), though the detected rates of jets are different in different categories of radio sources with the same instruments under similar working conditions. The different detected rates of jets can be reconciled with the unification model of active galaxy nuclei (AGNs) (Bridle \\& Perley \\cite{Bridle3}; Antonucci \\cite{Antonucci}). By studying jets, one can learn the physics of a jet itself, the ambient medium, and the central engine of the radio source. Many observational works have been dedicated to the studies of the structures of galaxies (e.g. Leahy et al. \\cite{Leahy} and references therein; Owen \\& Ledlow \\cite{Owen} and references therein; Xu et al. \\cite{Xu} and references therein, Morganti et al. \\cite{Morganti}; Condon et al. \\cite{Condon}; Parma et al. \\cite{Parma}; Carilli et al. \\cite{Carilli}), radio quasars and BL Lac objects (Lonsdale et al. \\cite{Lonsdale}; Murphy et al. \\cite{Murphy}; Bridle et al. \\cite{Bridle2}; Price et al. \\cite{Price}; Akujor et al. \\cite{Akujor}; Garrington et al. \\cite{Garrington}; Kollgaard et al. \\cite{Kollgaard}; Laurent-Muehleisen et al. \\cite{Laurent}) since the publication of the jet list given by Liu \\& Xie (\\cite{Liu}). These works describe a large number of kpc-scale jets in many radio sources. With Very Large Baseline Array (VLBA) and global Very Long Baseline Interferometer (VLBI), the structures of the center of many radio sources are investigated in detail and pc-scale jets (e.g. Spencer et al. \\cite{Spencer}; Taylor et al. \\cite{Taylor} and references therein; Dallacasa et al. \\cite{Dallacasa}; Fey \\& Charlot \\cite{Fey} and references therein) are revealed in many compact steep spectrum radio sources and flat-spectrum radio sources. A recent review is given by Zensus (\\cite{Zensus}) on the understanding of the properties of pc-scale radio jets in extra-galactic radio sources.  In their paper, Bridle and Perley (\\cite{Bridle3}) list 125 confirmed and 73 possible extra-galactic radio jets. In 1992, Paper I gave 276 radio sources with detected jets known to them prior to mid-December 1989. As many more detected extra-galactic radio jets have been reported in the literature since then, it is time to give a updated list of radio jets. To make the new list in this paper consistent with those reported in the literature, we adopted the morphological definition of a radio jet given by Bridle and Perley (\\cite{Bridle3}). In addition to the information on luminosities of jets, total sources and central cores, the length, sidedness and simple code of morphology were also given in this work. A detailed description of the catalogue will be given in Sect.~2. A simple discussion and conclusion will be given in Sect.~3. Jet properties in IR, optical, UV, and X-ray bands will be discussed later (Liu et al. \\cite{Liub}).  A Friedmann standard cosmology model is adopted and Hubble constant $H_0 = 100 Km \\cdot s^{-1} Mpc^{-1}$ and deceleration factor $q_0 = 0.5$ are used throughout this paper. ",
        "conclusions": "Jets are detected in the red-shift range from $ Z_{min} = 0.0012$ to $ Z_{max} = 3.886$. The closest radio source with detected jets is the nearest active galaxy Cen A (1327+427) at 3.7 Mpc. Jets are detected in QSOs in a red-shift range from $Z_{min} = 0.104$ to $ Z_{max} = 3.886$. The minimum value is very close the lower limit of red-shift defining a QSO in Hewitt and Burbidge (\\cite{Hewitt}). The maximum of the red-shift dramatically increases from 2.877 in Paper I to 3.886. Because of the importance of high red-shift radio galaxies in the study of the formation and evolution of galaxies, many very powerful instruments have been dedicated to the observations of radio galaxies with high redshift and consequently brought about the detection of a large number of jets. The highest red-shift for radio galaxies in Paper I is 0.574, which is much smaller than 3.395 given in Table 2. The ranges of total, jet, and core powers do not change a lot, and are comparable to those in Paper I.  Observations of compact steep-spectrum sources is another important subject in the theory of galaxy formation. Only pc-scale jets are detected in compact flat-spectrum and steep-spectrum sources. In the galaxy formation theory, compact steep-spectrum sources are at the early stages of the evolution of AGNs and the detected pc-scale jets are instrinsically small. On the other hand, compact flat-spectrum sources are active galaxies with a small view angle of the jet to the line of sight and the pc jets are the projections of the kpc jets.  We classify radio sources as radio galaxies, radio quasars, BL Lac objects, and Seyfert galaxies, based on the optical properties of the sources. Jets are commonly detected in all types of sources with different detection rates. Large-scale jets are detected in a large fractions of weak radio galaxies and extended radio quasars, although the fraction becomes lower in the lower range of power of core radio emission (Parma et al. \\cite{Parma}). This may result from the selection effect as a consequence of the correlation between opening angle and radio power (Bridle \\cite{Bridle1}; Parma et al. \\cite{Parma}; Paper I). Jets are detected in a smaller fraction of distant radio galaxies than in radio quasars with similar radio power and with instruments operating under similar parameters.  We gave in this paper a new list of 661 extra-galactic radio sources with detected jets. The list has two significant features: pc-scale jets are detected in many compact flat-spectrum and compact steep-spectrum sources and kpc-scale jets are detected in a large number of radio galaxies and radio quasars."
    },
    "0212/astro-ph0212194_arXiv.txt": {
        "abstract": "s{ We discuss a {\\it non-self-similar analytical model} capable of explaining the formation of accretion-powered galactic and extra-galactic jets. } \\vskip 0.5cm \\hrule \\noindent {\\bf Oral presentation, Session APT1. Published in the Proceedings of the IX th Marcel Grossmann Meeting.} \\hrule ",
        "introduction": "Though widely observed to be emanating from a variety of astrophysical sources, the underlying physical mechanism behind the formation of galactic and extragalactic jets is still enshrouded in a veil of mystery.  In addition, it has not been possible to calculate accurately the amount of barionic matter expelled in these events. Though progress in modelling the exact mechanism of jet formation has been severely handicapped by inadequacy in understanding the proper microphysics responsible jet formation, the basic ingradients of a theoretical model in attempt to explain the origin and formation of galactic/extragalactic jets seems to be well agreed upon \\cite{fer,mira,chak}. These are: firstly, a deep central potential well provided by an SMBH in case of jets from AGNs and QSOs and stellar mass BHs in case of jets from galactic Microquasars, and secondly, matter accreting onto these central compact objects to provide the energy for continuous throttling of the jets. Because of the fact that BHs don't have their own physical atmosphere from where matter could be ripped off as winds, outflows/jets in these cases {\\it have to be generated only from the accreting material}. This invariable association of accretion processes with jet formation reinforce the belief that for galactic and extra-galactic jet sources, inflow and outflow must be studied within the same framework. Also to be noted that while self-similar models are a valuable first step, they can  never be the full answer, and indeed any model which works equally well at all radii is fairly unsatisfactory to prove its viability. Thus the preferred model for jet formation must be one which is able to select the {\\it specific} region of jet formation.\\\\ ",
        "conclusions": ""
    },
    "0212/astro-ph0212531_arXiv.txt": {
        "abstract": "{  We have obtained HST-WFPC2 F555W and F814W photometry for 16 fields in the vicinity of the luminous nearby spiral galaxy M31, sampling the stellar content of the disk and the halo at different distances from the center, from $\\sim$ 20 to $\\sim$ 150 arcmin (i.e. $\\sim$ 4.5 to 35 kpc), down to limiting V and I magnitudes of $\\sim$ 27.  The Color-Magnitude diagrams (CMD) obtained for each field show the presence of complex stellar populations, including an intermediate age/young population and older populations with a wide range of metallicity. Those fields superposed on the disk of M31 generally show a blue plume of stars which we identify with main sequence members. According to this interpretation, we find that the star formation rate over the last 0.5 Gyr has varied dramatically with location in the disk.  The most evident feature of all the CMDs is a prominent Red Giant Branch (RGB) with a descending tip in the V band, characteristic of metallicity higher than 1/10 Solar. A red clump is clearly detected in all of the fields, and a weak blue horizontal branch is frequently present.  The metallicity distributions, obtained by comparison of the RGB stars with globular cluster templates, all show a long, albeit scantly populated, metal-poor tail and a main component peaking at [Fe/H] $\\sim$ -0.6. The most noteworthy characteristic of the abundance distributions is their overall similarity in all the sampled fields, covering a wide range of environments and galactocentric distances. Nevertheless, a few interesting differences and trends emerge from the general uniformity of the metallicity distributions. For example, the median $[Fe/H]$ shows a slight decrease with  distance along the minor axis (Y) up to $Y\\simeq 20 \\arcmin$, but the metallicity gradient completely disappears beyond this limit. Also, in some fields  a very metal-rich ([Fe/H] $\\ge$ -0.2) component is clearly present.  Whereas the fraction of metal-poor stars seems to be approximately constant (within few percent) in all fields, the fraction of metal-rich and, especially, very-metal-rich stars varies with position and seems to be more prominent in those fields superposed on the disk and/or with the presence of streams or substructures \\cite[e.g.][]{iba01}. This might indicate and possibly trace interaction effects with some companion, e.g. M32. ",
        "introduction": "As the nearest bright spiral and the most prominent member of the Local Group, M31 has played a central role in our evolving understanding of stellar populations. Most notably, Baade's (1944) identification of Population II in M31 initiated a true paradigm  shift  which in large part remains valid to the present day.  Therefore, studying in detail the stellar populations in M31 and comparing their properties to those of their Galactic counterparts is obviously very important to understand the formation and evolution history of these two galaxies. Our present knowledge already points to some interesting differences, supported by observational evidence \\citep[see] [and references therein]{hh02,fer02}, e.g.:  i) The globular cluster population in M31 is $\\sim$ 3 times larger than in the Milky Way, and on average more metal rich \\citep{mou86,du94,du01,barmby,per02};  ii) Also the abundance of the field population, in those few cases where it is derived from calibrated colors, seems generally higher than in the Milky Way, even at large galactocentric distances where the halo population dominates \\citep[cf.][and references  therein]{du01, fer02, rich96}. A similar result was found  for the halo stars in NGC 5128, with a striking difference compared to the Milky Way \\citep{hh00,hh02}. This has an impact on galaxy formation models, in particular those based on accretion of disrupted satellites, since dwarf galaxies are metal poor \\citep[][and references therein]{dco02}.  iii) \\citet{pri94} found that ``a  single de Vaucouleurs  law  luminosity profile can fit the spheroid of M31 from the inner bulge  all the way out to the halo''. However, clear evidence of ``disturbances'' such as spatial density and metallicity variations has been found, that can be interpreted as streams and remnants of tidal interactions  \\citep{iba01, fer02}.  The available data, therefore, would point toward interaction/merger events as non negligible factors in the building of the M31 halo. Therefore, very deep and detailed studies of the stellar  populations and abundance distributions in the M31 halo and outer disk are needed, as they can play a fundamental r\\^ole in explaining the formation of the M31 halo.   The advent of new technology detectors and telescopes, in particular the Hubble Space Telescope ($HST$), has boosted a new generation of studies, providing a much deeper insight in the understanding of the stellar content and evolutionary  history of this galaxy. In line with these recent efforts \\cite[e.g., see][and references therein]{du94,mo94,cou95,rich96,hof96,ho98, guha00,fj01,du01,ste,svd,fer02,wil02}, we present here the first results of a large systematic survey of the stellar population in the disk and halo of M31 performed with the Wide Field and Planetary Camera-2 (WFPC2) on board $HST$.  The $HST$ program GO-6671  (PI R. M. Rich) was aimed primarily at imaging with the WFPC2 a sample of bright globular clusters in M31 spanning a range in metallicity.  The globular cluster target is placed on the Planetary Camera (PC) field, while the adjacent halo or disk field population of M31 falls on the 3  Wide Field Camera (WFC) chips. Adding to these  7 more fields adjacent to globular clusters, whose images are taken from the $HST$ archive, brings a total of 16 deep fields imaged by $HST$, over a distance of about 4.5 to 35 kpc from the nucleus. More M31 fields are available in the $HST$-archive, but we analyze here only those obtained in a strictly homogeneous way (i.e. same camera and filters, similar exposures).  Photometry of the clusters is considered in a separate paper \\citep{rich03}. In this paper we consider only photometry of the fields imaged by the WFC (Wide Field Camera) chips. Our photometry generally reaches to $\\sim 1 $ mag fainter than the horizontal branch.  Our primary aim is to derive the abundance distribution of the M31 fields assuming that the population is old and globular cluster-like. This assumption is well justified by deep photometry \\cite[e.g.][]{rich96} that shows the subgiant luminosity function to approximately match that of the old  globular cluster 47 Tuc. The abundance distribution of the M31 field stellar population is then derived by interpolating between empirical template globular cluster RGB ridge lines, a technique first applied by \\citet{mou86}  and later used in most of the subsequent studies, though with significant differences.   The paper is organized as follows.  Section 2 discusses the properties of our field locations within the Andromeda galaxy. Section 3 discusses our observations and data analysis. Section 4 presents the individual color-magnitude diagrams, including an analysis of the young main sequence (blue plume) stellar population.  Section 5 describes our method for deriving the abundance distribution of the old stellar population. Section 6 reports the abundance distributions, while Section 7 considers the error budget (sensitivity to parameters such as reddening, distance modulus and abundance scale). We compare our abundance distributions with those in the literature in Section 8, while Section 9 considers how the detailed abundance distributions vary with position and environment. We summarize our results in Section 10. ",
        "conclusions": "We have analysed HST-WFPC2 images of 16 fields in M31 at a wide range of distances from the galactic center and plane. These fields have been observed in the F555W (V) and F814W (I) filters and reach a roughly uniform depth.  The color-magnitude diagrams of these fields are generally dominated by a red giant branch and a populous red clump. The RGB has a strong descending tip, indicative of a metal rich population. In the outermost fields a blue Horizontal Branch is also detected. When present, the blue HB represents about 15\\%\\ of the total HB population, roughly estimated from  the fraction of metal poor giants.  In fields superposed on the disk, a blue plume of main sequence stars is also identified. We report a simple analysis of the blue plume population, which shows that the star formation history of the disk has been spatially inhomogenous during the last 0.5 - 1 Gyr.  We have obtained photometric metallicity distributions from RGB stars by interpolation on a grid of empirical globular cluster RGB templates.  The most robust result of the present $HST$ survey is the evidence that the metal-rich population, with  $[Fe/H] \\sim -0.6$, is the major component of the stellar mix everywhere in the $10 \\le R_{arcmin} \\le 130$ range. A  minor metal-poor component (with $[Fe/H]<-1.4$) is also ubiquitous. The old stellar population is remarkably uniform across the disk, outer halo fields, and in the proposed tidal stream.  This uniformity is the dominant feature of the old stellar population.  This basic result is not new \\cite[see, e.g.,][and references therein]{du01}, but coupled with data shown by the very wide mapping carried out by \\citet{fer02}, fixes on solid grounds the conclusion that the stellar population of the M31 spheroid differs substantially from that of the Milky Way and is almost an order of magnitude more metal rich, on average. In this respect, we confirm also the finding of \\citet{rich96}, that the metal rich population is present even at 150 arcmin ($\\sim$ 35 kpc) from the nucleus, at the location of the globular cluster G1.  An obvious explanation for the origin of the metal rich halo of M31 is that the population formed in connection with the outflow of metal-enriched winds, perhaps associated with the formation of the spheroid.  Metal rich populations in halos are clearly widespread \\cite[e.g. NGC 5128; ][]{sor97,hh99,hh02}. There is evidence for the outflow of metal enriched winds in Lyman break galaxies, and for near-solar abundances in these galaxies, at high redshift \\citep{ste96}. The stars formed from this material might well live to comprise the population II halo.   The outflow of enriched material will not affect the shape of the abundance distribution (the classic Simple Model form is retained) but the yield (mean abundance of the whole galaxy) does decline \\citep{hart76}.  The increasing fraction of the \"Metal-Poor Population\" (i.e. with $[Fe/H]< -0.8$) with distance from the plane Y suggests a dissipational origin for the regions that are closer to the galactic plane.   This evidence is compatible with the hypothesis that a \"homogeneous\"  formation event may have occurred in the very early stages of the  M31 history.  On the other hand, the image of the M31 halo illustrated by \\citet{iba01} and \\citet{fer02} shows (for the metal-rich stars taken as a {\\em global} population) a quite flattened, disk-shaped distribution, associated with the outskirts of the inner disk. It is possible thus that the metal rich population might be associated with a proto disk or flattened halo.  The large number of stars involved would strongly favor the notion that we study the {\\it old } disk, and not just some trace population such as the thick disk.  So we have this alternative scenario, where the dominant metal-rich  population might be associated with a meta-disk extending perhaps more the 20 kpc, while the metal-poor population might be  associated  with the spheroid. The two populations -- and the corresponding structures -- would have different evolutionary histories. The old stellar disk would be much larger than previously believed, extending out to a few degrees from the nucleus of M31 and dominating the stellar mix also along lines of sight very distant from the center of the galaxy \\citep{fj01,fer02}. It is also possible that the inner disk and bulge had common, or closely tied, formation histories.  In addition to this, the sub-structures (``streams'' and ``clumps'') clearly detected by \\citet{iba01,fer02} and supported by our results might be related to the phenomena of merging or interaction with satellites. In this case the metal rich component might be connected with the tidal disintegration of a companion (\\citet{fer02}), a likely candidate being M32, perhaps involving a good deal of its mass \\citep{bek02,cho02}.  New data are now required to choose among the scenarios presented here.  Deeper HST imaging, some of which is in progress, may constrain the actual age distribution of stars in some regions of the halo. Large scale spectroscopic surveys of halo stars have just begun, and have the potential to determine what fraction of the M31 halo might have originated from the tidal disruption of satellites.  Our first round of large-scale HST imaging gives evidence that the old ``halo'' stellar population is more metal rich than that of the Milky Way, with surprisingly little variation in the properties of the old stellar population, even for fields ranging up to galactocentric distances exceeding 30 kpc. The challenge now will be assemble a data set for the Galaxy and M31 powerful enough to constrain their origins."
    },
    "0212/astro-ph0212257_arXiv.txt": {
        "abstract": "{The new World Wide Web site \"GOLDMine\" (Galaxy On Line Database Milano Network) (http://goldmine.mib.infn.it) contains a multiwavelength data-base of an optically selected sample of 3267 galaxies in the  Virgo cluster and in the  Coma Supercluster. It is designed for professional astronomers who wish to find data and images for these galaxies. Data, gathered in 15 years of observational campaigns by the authors or taken from the literature include  general parameters (catalogue names, celestial coordinates, morphological type, recessional velocity etc.); multiwavelength  continuum photometry (total UV, U, B, V, J, H, K, FIR and radio magnitudes/flux densities); line photometry (HI, H$_2$, H$\\alpha$);  dynamical parameters (rotational velocity from the HI and H$\\alpha$ lines, velocity dispersion) and structural parameters (light concentration index, effective radius and brightness, asymptotic magnitude) in the optical (B and V) and Near Infrared (H or K) bands.\\\\ Images include finding charts, optical (B and V), H$\\alpha$, Near Infrared (H and/or K) and true color RGB frames (when available). Radial light profiles obtained from the B, V, H or K band images are also available. Integrated optical spectra along with broad Spectral Energy Distributions (SED) from the UV to the radio domain are given. All images can be obtained in JPG format, but the original (reduced) FITS images can be downloaded as well. The database will be updated regularly and will be extended to other local clusters and superclusters. Astronomers who wish to have their images included in GOLDMine are strongly encouraged to send us their material. } ",
        "introduction": "Galaxies come in many different forms and sizes, but they can be broadly divided into two main species: {\\it Spirals}, with a flattened, disk-like shape, blue colors, much gas and dust, and a widespread star formation activity that results in the presence within them of many young stars and {\\it Elliptical}, with a spheroidal shape, red colors, little or no gas and dust, and no star formation activity, thus containing exclusively old stars. While these differences are well documented, and have been extensively studied over the last 70 years, astronomers are not yet able to convincingly explain the origin of such a diversity. In fact one of today's major open issues for astronomy is to obtain a plausible reconstruction of the processes of galaxy formation and evolution. The task is of course enormously complicated by the sheer disparity between cosmological and human time-scales, that makes it impossible for astronomers to witness galaxy evolution directly. One possible way to overcome this problem is to take advantage of the ''time machine'' effect provided by the finite speed of light. Observing today galaxies at different distances means observing them at different epochs in the history of the Universe, and thus with different ages.  Unfortunately distances must be very large before age differences become significant, and therefore unveiling the evolution of galaxies with direct observations of very distant objects is becoming a reality only today, with the advent of the 10m class telescopes.  Galaxies at progressively larger look-back time (redshift), thus in a younger evolutionary stage, will be observed in the near future, back to a small fraction of the present age of the Universe. This will provide us with a sequence of \"fossil\" galaxies, eventually disclosing the secret of their evolution, much as fossil organisms guided entomologists tracing back the evolution of species.  However, 15 years ago, when we initiated this project, a direct observational approach to the problem was unconceivable, because of the lack of proper telescopes and instrumentation. The only available option was to attack the evolutionary enigma starting from the study of evolved systems, i.e. the nearby galaxies. This approach is based on the conviction that \"adult\" galaxies still preserve some memory of their past, and that adequate observations would eventually help to disclose it.\\\\ With this purpouse we started in 1985 an observational campaign aimed at providing the phenomenology of local galaxies in the widest possible frequency range. We took observations and collected data from the literature from the 2000 \\AA (UV) range to the centimetric radio domain, spending a large effort in making the literature data as homogeneous as possible with our own data. GOLDMine is designed to provide access to this massive data-set on local galaxies trough the World Wide Web. Numeric parameters as well as the (reduced) scientific FITS images can be downloaded. ",
        "conclusions": "We undertook a 15 years research project focused on the properties of galaxies in the local universe which brought some new insight on their phenomenology and possibly will help disclosing their origin and evolution (see Gavazzi 1993; Gavazzi \\& Scodeggio 1996; Gavazzi et al. 1996a, 2002; Boselli et al. 2001). \\begin{itemize} \\item{A complete, optically selected sample of 3267 galaxies representative of all morphological types and luminosities was extensively surveyed through various observational windows from the UV to the centimetric radio.} \\item{The surveyed regions include the Virgo cluster and the Coma supercluster.} \\item{Numeric parameters and images are made available to the international community through the World Wide Web site http://goldmine.mib.infn.it.} \\end{itemize} If your research benefits from the use of GOLDMine, we would appreciate the following acknowledgement in your paper: \"This research has made use of the GOLDMine Database, operated by the Universita' degli Studi di Milano- Bicocca\". Please also cite the present paper."
    },
    "0212/astro-ph0212061_arXiv.txt": {
        "abstract": "If Ultra-High Energy Cosmic Rays (UHECRs) are originated from nearby galaxies, modeling of the distribution of nearby galaxies is important to an accurate estimate the source number density of UHECRs. We investigate uncertainty of the selection function of the Optical Redshift Survey (ORS), which we used to construct a source model of UHECRs. The investigation is based on a comparison of numbe counts of ORS galaxies with those of the spectroscopic sample of the Sloan Digital Sky Survey (SDSS) Early Data Release (EDR). We carefully count galaxies in the same absolute magnitude bin from the two samples. We find a slight systematic overestimate of the ORS counts outside 5000 km s$^{-1}$ by about a factor of 2. We revise the selection function of the ORS assuming that the SDSS counts are correct. Our revision is based on the absorption given in the ORS catalog as well as that computed from Schlegel et al. (1998), which is systematically larger than the former by $A_B \\sim 0.1$ mag in the region of low absorption. It is found that introduction of Schlegel et al.'s absorption changes one of the parameters of the ORS selection function by more than $10 \\%$. The revision should be taken into account in the future analysis of the source number density of UHECRs based on the ORS. Using the revised selection function, we determine the global structure of the Local Supercluster (LSC) with a source model of UHECRs, that is, a number-density model consisting of a uniform spherical halo and an exponential disk. We find that the revision is insignificant in terms of the structure of the LSC. However, the revised selection function will be useful to other studies such as peculiar velocity and correlation function. ",
        "introduction": "\\label{data} \\indent  \\subsection{THE ORS SAMPLE} \\label{ors} \\indent  The ORS is drawn from three catalogs, the European Southern Observatory Galaxy Catalogue (ESO) covering the southern sky ($\\delta \\leq -17.5^{\\circ}$), the Uppsala Galaxy Catalogue (UGC) in the northern sky ($\\delta  \\geq -2.5^{\\circ}$), and the Extension to the Southern Galaxy Catalogue (ESGC) in the strip between the ESO and UGC region ($-17.5^{\\circ} \\leq \\delta  \\leq -2.5^{\\circ}$). The last catalogue contains no magnitudes other than very poor ones from eye inspection. The sample selection and galaxy distribution are described in detail by Santiago et al. (1995). This survey contains two subcatalogs, one complete to a B magnitude of 14.5 (ORS-m) and the other complete to a B major axis diameter of $1.'9$ (ORS-d). In this study we calibrate the ORS selection function using the SDSS EDR which is a magnitude-limited survey. Accordingly, we use the magnitude-limited subcatalog, i.e., ORS-m, which contain 5716 galaxies. The subcatalog ORS-m itself is also composed of two subsamples, ESO-m (2437 galaxies) and UGC-m (3279 galaxies). We further restrict ourselves to galaxies within 8000 km s$^{-1}$ and with $M_B - 5$log$_{10}h < -0.5 - 5$log$_{10}v$, which corresponds to the apparent magnitude limit of $m_B=14.5$, where $h$ is the dimensionless Hubble constant $h=H_0/(100$ km s$^{-1}$ Mpc$^{-1})$ and $v$ is the recession velocity. Very faint nearby galaxies with $M_B>-14.0$ are excluded. Galaxies in the zone of avoidance ($|b|<20^{\\circ}$) are also excluded. Our final ORS sample contains 5178 galaxies with 2346 galaxies in the ESO-m and 2832 galaxies in the UGC-m.  Figure~\\ref{fig1} shows the sky distribution of our ORS galaxies in the equatorial coordinate. Figure~\\ref{fig2} shows the absolute magnitude plotted as a function of the recession velocity. Throughout the paper, we assume that there are no departures from uniform Hubble expansion for simplicity. We use $h$ = 1 when we do not specify $h$ explicitly. An alternative value for $h$ changes all absolute magnitudes by 5log$_{10} h$.     In order to compare the galaxy number counts of the ORS with those of the SDSS EDR, a selection function has to be introduced. This selection function of the ORS is derived for several subsamples in Santiago et al.(1996) using galaxies to 8000 km s$^{-1}$. They assumed that the selection function is expressed as \\begin{equation} \\phi_{\\rm ORS}(r) = \\left \\{ \\begin{array}{ll} \\left( \\frac{r}{r_{s}}\\right) ^{-2a} \\left( \\frac{r_{*}^2 + r^2}{r_{*}^2 + r_{s}^2} \\right) ^{-b} &\\quad r > r_{s} \\\\ 1 &\\quad r \\leq r_{s} \\end{array} \\right., \\end{equation} where $r$ represents distance from our Galaxy, and they set $r_{s}$ to be 500 km s$^{-1}$. Values of the parameters determined by Santiago et al.(1996) are $a$ = 0.40, $b$ = 6.25, $r_{*}$ = 11100 (km s$^{-1}$) for ESO-m and $a$ = 0.36, $b$ = 9.38, $r_{*}$ = 14060 (km s$^{-1}$) for UGC-m. The flux limit $m_B = 14.5$ corresponds to $M_B = -14$ at $v = 500$ km s$^{-1}$, where the selection function is unity. In Figure~\\ref{fig3} we show the selection function of the ORS given by Santiago et al.(1996) as a function of recession velocity together with that of the SDSS (explained in the next subsection).    \\indent \\subsection{THE SDSS SAMPLE} \\label{sdss} \\indent   \\subsubsection{Completeness} \\label{complete} \\indent  The SDSS is the imaging survey in five bands ($u', g', r', i', z'$) followed by spectroscopy. The spectroscopic sample of the SDSS EDR contains 1504 galaxies with measured recession velocity $v<8000$ km s$^{-1}$. The region covered by the spectroscopy of the SDSS EDR is shown in Figure~\\ref{fig4}. In the following, we refer to the region at $\\delta>50^{\\circ}$ as region 1, two regions in the equatorial stripe as regions 2 and 3. All the data in these three regions are collectively referred to as the ALL region.    At the brightest end, it is suspected that problems in the deblender and/or possible saturation limit the completeness of the SDSS. There are 130 ORS galaxies which fall in the ALL region. We tried to match these 130 galaxies with the SDSS catalogs. We found that 72 ORS galaxies out of the 130 could be matched with the SDSS spectroscopic objects and that 43 could be matched with the photometric objects (but not with the spectroscopic objects). Remaining 15 galaxies are not included either in the spectroscopic catalog nor in the photometric catalog. Among them, 10 suffer from problems in the deblender, 1 is too bright ($\\sim$ 9 mag), and remaining 4 galaxies are missed by some unknown reason. Thus the completeness of the SDSS spectroscopic catalog is quite low, $72/130=55\\%$, at the ORS limiting magnitude. However, if supplemented by the photometric catalog, it increases to $115/130=88\\%$. The SDSS spectroscopic catalog have a very high completeness for galaxies including very faint ones (above $99\\%$ for $r'<18$) (Strauss et al. 2002). We add the 43 SDSS photometric objects in the above 1504 galaxies assigning them the recession velocities given in the ORS. Thus, our SDSS sample now includes 1547 galaxies. We extract our final sample from these 1547 galaxies to be compared with the ORS sample including the 130 galaxies.  \\indent \\subsubsection{Conversion from $m_B$ to $g'$ magnitude} \\label{conversion} \\indent  The ORS B magnitude is based on the photographic plates and its bandpass is similar to that of Couch-Newell B$_J$ (Couch and Newell 1980). We use only SDSS $g'$-band magnitude in this study because it is close to the B$_J$ magnitude. In order to carefully count galaxies in the same magnitude range, we have to know magnitude difference between $m_B$ and $g'$. For this purpose, we made a direct comparison of the magnitudes of the matched 72 plus 43 galaxies between the two surveys. However, variance of the magnitude offset was very large ($g'-m_B \\sim 1.5$ at the maximum). At such bright end (the ORS-m is complete to $m_B=14.5$ mag), it is suspected that problems in the deblender limit the accuracy of the SDSS magnitude. Accordingly, we derive the magnitude offset in the method explained below.    There are three differences of the magnitude system between the SDSS and the ORS. First, response functions of these two photometric band systems differ from each other. Magnitude offset due to this difference can be computed from galaxy colors in various photometric band systems given in Fukugita et al.(1995) for typical morphology types grouped into E, S0, Sa-Sb, Sb-Sc, Sc-Sd, and Sdm-Im. The ORS gives the morphological type index to individual galaxy, whereas the SDSS does not. However, Strateva et al.(2001) showed that galaxies have a bimodal $u^*-r^*$ color distribution corresponding to early (E, S0, Sa) and late (Sb, Sc, Irr) morphological types and that the two types can be clearly separated by a $u^*-r^*$ color cut of 2.22 independent of magnitude. Accordingly, we compute the magnitude offset for early and late morphological types and use them throughout the paper. Averaging over morphology types of galaxies in our ORS sample, we adopt $g'-B_J=-0.35$ and $-0.24$ for early and late type galaxies, respectively.   Second, there is a difference about the method for measuring the flux. In order to measure a constant fraction of the total light, independent of Galactic absorption and cosmological dimming, the SDSS has adopted a modified form of the Petrosian (1976) magnitude. Magnitudes are measured within the aperture of twice the Petrosian radius, which is insensitive to absorption or cosmological dimming. The detailed explanations are presented in Blanton et al.(2001), Yasuda et al.(2001), and Stoughton et al.(2002). On the other hand, the ORS is based on an isophotal magnitude. The relationship between the SDSS Petrosian magnitudes and isophotal magnitudes is examined in Blanton et al.(2001). Difference between the two magnitudes $\\Delta m = m_{petro}-m_{iso}$ ranges from $-0.2$ to $0.1$ according to several choices of the isophotal limit in case of $r'$ band. It is, however, not known exactly what value we should take for $\\Delta m$ in the $g'$ band. Here we tentatively take $\\Delta m = -0.1$ with an estimated uncertainty of $\\pm 0.2$.   Finally, SDSS magnitudes are not the traditional logarithmic magnitude but inverse hyperbolic sine magnitude, which is described in detail by Lupton et al.(1999). However, these magnitudes differ by less than 1 \\% in flux above $g'=22.60$. We neglect this difference throughout the paper.   Taking account of these three differences between the SDSS and the ORS magnitudes, we take the total amount of the magnitude offset as $g'-B_J+\\Delta m = -0.45 \\pm 0.2$ and $-0.34 \\pm 0.2$ for early and late type galaxies.  \\indent \\subsubsection{Selection of galaxies from our SDSS sample} \\label{selection} \\indent   We apply the following selection criteria to our 1547 SDSS galaxies with $v \\leq 8000$ km s$^{-1}$. The limiting magnitude is set at $g'=17.65$, which was adopted in the luminosity function study by Blanton et al.(2001). This corresponds to the limit in the absolute magnitude of $M_{g'} = 2.65-5 $log$_{10}v$. The limiting magnitude of the ORS is $M_B=-14.0$, which corresponds to $M_{g'}=-14.45$ and $-14.34$ for early and late type SDSS galaxies. Accordingly we also exclude galaxies fainter than this limiting magnitude. The absolute $g'$ band magnitude plotted as a function of the recession velocity is shown in Figure~\\ref{fig5}. The sudden change in the density of galaxies which occurs at around 5000 km s$^{-1}$ reflects a real large scale structure. Among the 1547 galaxies we find a tiny fraction ($\\sim$ 3\\%) of galaxies with $g' \\gg 17.65$, which may be assigned unrealistically faint absolute magnitudes because of some problem in deblending. There is concern that these galaxies are actually more luminous than $g' = 17.65$. They are not included either in the sample or in Figure~\\ref{fig5} because of the above cut in the absolute magnitude. However, rejection or inclusion of this tiny fraction of erroneous data does not affect our results. Our final SDSS sample contains 1078 galaxies. Figure~\\ref{fig6} shows the distribution of galaxies in our SDSS sample and ORS galaxies in the EDR region in the equatorial coordinate. The numbers of observed galaxies in our SDSS and ORS sample within each region are tabulated in Table~\\ref{num}. Note that the whole region covered by the SDSS sample is located in the area covered by UGC-m (See Figure~\\ref{fig1}).    \\indent \\subsubsection{Selection function} \\label{sf} \\indent   The selection function of the SDSS can be computed by using the luminosity function (Blanton et al. 2001). The luminosity function can be expressed by the Schechter function as \\begin{equation} \\Phi(M) \\propto 10^{-0.4(M-M_{*})(\\alpha + 1)} \\times \\exp \\left[10^{-0.4(M-M_{*})}\\right] \\quad, \\end{equation} where M is the absolute magnitude. The parameters for $g'$ band determined by Blanton et al.(2001) are $M_{*} = -20.04 \\pm 0.04$ mag, $\\alpha = -1.26 \\pm 0.05$ for $\\Omega_{m} = 0.3$ , $\\Omega_{\\Lambda} = 0.7$ cosmology. When we restrict our analysis to galaxies with absolute magnitudes $M < M_{0}$, the selection function is given by \\begin{equation} \\phi_{\\rm SDSS}(r) = \\frac{\\int_{-\\infty}^{M_{\\rm max}(r)} \\Phi(M) {\\mathrm d} M}{\\int_{-\\infty}^{M_{0}} \\Phi(M) {\\mathrm d} M}, \\label{selecsdss} \\end{equation} where $M_{\\rm max}(r) =  $min$(M_{0},M_{\\rm lim}(r))$ and $M_{\\rm lim}(r)$ is the limiting magnitude for a galaxy at distance $r$ to be included into the sample. We set $M_{0}$ to be $-14.45$ and $-14.34$ for early and late type galaxies which are the absolute magnitudes in $g'$ band corresponding to $M_B = -14$. The SDSS selection function obtained in this way is shown in Figure~\\ref{fig3}.  \\indent ",
        "conclusions": "\\label{summary} \\indent   In this paper, we investigate the uncertainty of the selection function of the ORS, which we used as sources of UHECRs, by comparing the galaxy number counts of ORS galaxies with those of SDSS EDR galaxies as a function of recession velocity. We carefully count galaxies from the two data samples so that these galaxies are in the same absolute luminosity range. In all the regions covered by the SDSS EDR, we find that the ORS predicts systematically more galaxies than the SDSS by a factor of $\\sim$ 2 outside 5000 km s$^{-1}$. We revise the selection function of the ORS so that the predictions of the two surveys become consistent. It is also found that replacing the original absorption given in the ORS catalog with that computed from Schlegel et al.(1998) changes the $b$ parameter of the ORS selection function by more than 10 $\\%$. The revision should be taken into account in the future analysis of the source number density of UHECRs based on the ORS.    Using this revised selection function of the ORS, we determine the spatial structure of the LSC with a model consisting of a uniform spherical halo and an exponential disk. The model parameters are derived for galaxies in three ranges of different absolute luminosities, $M_B < -14$, $M_B < -16$ and $M_B < -18$. The most probable values of the scale length of the disk in $z$ direction, $\\beta$, are almost the same for the three luminosity ranges while the radial scale length of the disk, $\\alpha$, and the density of the halo, $n_1$, are larger for $M_B < -18$ than for $M_B < -14$ and $-16$. This implies that giant galaxies $(M_B < -18)$ in the disk are more extended radially than dwarf galaxies $(M_B > -16)$ and that giant galaxies show a smaller local overdensity than dwarf galaxies.    Let us consider the effect of this difference of the distribution of galaxies in the LSC on the distribution of UHECR arrival directions. Observationally, it is reported that there is no significant large-scale anisotropy in the arrival direction distribution of UHECRs (Takeda et al. 1999). If all the galaxies, both giants and dwarfs, are sources of UHECRs, arrival directions could be more anisotropic than the case that only luminous giant galaxies are responsible for UHECRs. The observed lack of anisotropy may favor the latter case. Provided that UHECRs are generated in star forming regions, their sources are expected to trace the distribution of late-type galaxies rather than early-type galaxies. It is known that there is strong spatial segregation of the Hubble types in the Virgo cluster (Binggeli et al. 1987). Early-type galaxies are more concentrated towards the cluster center than late types. Distribution of galaxies of different morphological types in the LSC is also an important factor to be taken into account in the analysis of arrival direction of UHECRs.   We restricted ourselves in this paper to the study of galaxies as the sources of UHECRs. The revision of the ORS selection function made in this study is found to be insignificant in the determination of the global structure of the LSC. The revised selection function obtained in this study will, however, be useful to other studies such as peculiar velocity and correlation function. The SDSS will produce imaging and spectroscopic surveys in five bands ($u',g',r',i',z'$) over $\\pi$ steragians in the northern Galactic cap, and will collect spectra of approximately 1,000,000 galaxies (down to $r'_{lim} \\sim 17.65$). We have used in this paper the SDSS EDR which covers only less than $1\\%$ of all sky. As data obtained by the SDSS accumulate, we will be able to revise the selection function of several surveys with better accuracy, and also to determine the spatial structure of the LSC using only the SDSS galaxies."
    },
    "0212/astro-ph0212311_arXiv.txt": {
        "abstract": "s{ By studying the bispectrum of the galaxy distribution in the 2dF galaxy redhsift survey (2dFGRS), we have shown that 2dFGRS galaxies are unbiased tracers of the mass distribution. This allows us the break the degeneracy intrinsic to power spectrum studies, between the matter density parameter $\\Omega_m$ and the bias parameter $b$, and to obtain an accurate  measurement of $\\Omega_m$: $\\Omega_{m,z_{eff}}=0.27 \\pm 0.06$,  a measurement obtained from  the 2dFGRS  alone, independently from other data sets. This result has to be interpreted as $\\Omega_m$ at the effective redshift of the survey $z_{eff}=0.17$ Extrapolated at $z=0$ we obtain $\\Omega_m=0.2 \\pm 0.06$ for the $\\Lambda$CDM model. This constraint on the matter density parameter when combined with cosmic microwave background constraints on the flatness of the Universe, show firm evidence that the Universe is dominated by a vacuum energy component.} ",
        "introduction": " ",
        "conclusions": "We have demonstrated that the optically selected galaxies from the  2dFGRS trace the dark matter density extremely well on large scales. This allows us to break the degeneracy, intrinsic to power spectrum studies,  between gravity and biasing. We obtain a measurement of the matter density parameter at $z_{eff}=0.17$ of $\\Omega_{m,z_{eff}}=0.27 \\pm 0.06$. The extrapolation of this quantity at $z=0$ is model dependent, however the constraints of cosmological parameters obtained from 2dFGRS, Cosmic microwave background and Supernovae agree remarkably well. Thus we can cosnculd that we  are converging towards a cosmological model where the Universe is flat, is dominated by a vacuum energy component ($\\Lambda=0.8 \\pm 0.08$) and the matter density is $\\Omega_m=0.2 =\\pm 0.06$.      \\begin{figure} \\psfig{figure=combo.eps,height=2.5in} \\caption{LEFT: the 2dF galaxy redshift survey. RIGHT:1,2 and 3 sigma confidence contours in the $\\Omega_m$, $\\Omega_{\\Lambda}$ plane: green are the supernovae 1A consraints (Perlmutter et al 1999), red are the  CMB constraints from Boomerang (de Bernardis et al. 2000), blue are the LSS constraints obtained by the present work (Verde et al 2002). The solid lines show the joint likelihood contours of CMB +LSS data sets. It is truly remarkable that different measuerement of cosmological paraneters that use different observables and  are subject to different systematics agree so well. \\label{fig:cmblsssn}} \\end{figure}"
    },
    "0212/astro-ph0212133_arXiv.txt": {
        "abstract": "{  With the aid of simple numerical models, we discuss a particular aspect of the interaction between stellar orbital periods inside elliptical galaxies (Es) and the parent cluster tidal field (CTF), i.e., the possibility that {\\it collisionless stellar evaporation} from Es is an effective mechanism for the production of the recently discovered intracluster stellar populations (ISP). These very preliminary investigations, based on idealized galaxy density profiles (such as Ferrers density distributions) show that, over an Hubble time, the amount of stars lost by a representative galaxy may sum up to the 10\\% of the initial galaxy mass, a fraction in interesting agreement with observational data.  The effectiveness of this mechanism is due to the fact that the galaxy oscillation periods near its equilibrium configurations in the CTF are of the same order of stellar orbital times in the external galaxy regions. ",
        "introduction": "Observational evidences of an Intracluster Stellar Population (ISP) are mainly based on the identification of {\\it intergalactic} planetary nebulae and red giant branch stars (see, e.g., Theuns \\& Warren 1996, M\\`endez et al. 1998, Feldmeier et al. 1998, Arnaboldi et al. 2002, Durrell et al. 2002).  Overall, the data suggest that approximately 10\\% (or even more) of the stellar mass of the cluster is contributed by the ISP (see, e.g., Ferguson \\& Tanvir 1998).  The usual scenario assumed to explain the finding above is that gravitational interactions between galaxies in cluster and/or interactions between the galaxies with the tidal gravitational field of the parent cluster (CTF), lead to a substantial stripping of stars from the galaxies themselves.  Here, supported by a curious coincidence, namely by the fact that {\\it the characteristic times of oscillation of a galaxy around its equilibrium position in the CTF are of the same order of magnitude of the stellar orbital periods in the external part of the galaxy itself}, we suggest that an additional effect is at work, i.e. we discuss about the possible ``resonant'' interaction between stellar orbits inside the galaxies and the CTF.  In fact, based on the observational evidence that the major axis of cluster Es seems to be preferentially oriented toward the cluster center, N-body simulations showed that galaxies tend to align reacting to the CTF as rigid (Ciotti \\& Dutta 1994).  By assuming this idealized scenario, a stability analysis then shows that this configuration is of equilibrium, and allows to calculate the oscillation periods in the linearized regime (Ciotti \\& Giampieri 1998, hereafter CG98).  Prompted by these observational and theoretical considerations, in our numerical explorations we assume that the galaxy is a triaxial ellipsoid with its center of mass at rest at the center of a triaxial cluster. The considerably more complicate case of a galaxy with the center of mass in rotation around the center of a spherical cluster will be discussed elsewhere (Ciotti \\& Muccione 2003, hereafter CM03). ",
        "conclusions": "In Fig. 1 we show the results of one of our preliminary simulations for a galaxy model with $n=1$, $\\Mg =10^{10}\\,\\msun$, semi-mayor axis $\\alphau =20$ kpc, flattenings $\\epsilon =0.2$, $\\eta=0.4$, and maximum oscillation angles equals to 0.2 rad. The cluster parameters are $\\acun=\\sigvn =1$, $\\mu=0.2$, $\\nu=0.4$, and the total number of explored orbits is $\\Ntot=10^5$. In the ordinate axis we plot the number of {\\it escapers} as a function of their homeoidal radius at $t=0$.  Note how the zone of maximum escape is near $m\\simeq 0.8$, and how the total number of escapers is a significant fraction of the total number of explored orbits (actually, with the galaxy and cluster parameters adopted in the simulations, $N_{\\rm esc}\\simeq 0.1\\Ntot$). This number is a somewhat upper limit of the expected number in more realistic simulations: in fact, 1) as described in Section 4, we adopted a very weak escape criterium, and 2), in more realistic galaxy models the density decrease in the outer parts is stronger than in Ferrers models, and so we expect that a smaller number of stars will be significantly affected by the CTF. Both points are addressed and discussed elsewhere, in MC03 in a preliminary and qualitative way for an Hernquist model at the center of a cluster, while in CM03 we present the results of a systematic exploration of the parameter space for untruncated galaxies both at the center and with the center of mass in circular orbit in a spherical cluster.  In any case, the simple model here explored looks promising, with a number of escapers in nice agreement with the observational estimates."
    },
    "0212/astro-ph0212419_arXiv.txt": {
        "abstract": "The DASI discovery of CMB polarization % has opened a new chapter in cosmology. Most of the useful information about inflationary gravitational waves and reionization is on large angular scales where Galactic foreground contamination is the worst, so a key challenge is to model, quantify and remove polarized foregrounds. We use the POLAR experiment, COBE/DMR and radio surveys to provide the strongest limits to date on the $TE$ cross power spectrum of the CMB on large angular scales and to quantify the polarized synchrotron radiation, which is likely to be the most challenging polarized contaminant for the MAP satellite. We find that the synchrotron $E$- and $B$-contributions are equal to within 10\\% from $408-820$MHz with a hint of $E$-domination at higher frequencies. We quantify Faraday Rotation \\& Depolarization effects in the two-dimensional $(\\ell,\\nu)$-plane and show that they cause the synchrotron polarization percentage to drop both towards lower frequencies and towards lower multipoles. \\bigskip ",
        "introduction": "The recent discovery of cosmic microwave background (CMB) polarization by the DASI experiment \\cite{kovac02} has opened a new chapter in cosmology -- see \\fig{summary}. Although CMB polarization on degree scales and below can sharpen cosmological constraints and provide important cross-checks \\cite{ZSS97,Eisenstein98}, the potential for the most dramatic improvements lies on the largest angular scales where it provides a unique probe of the reionization epoch and primordial gravitational waves. For instance, forecasts \\cite{foregpars,knox02} indicate that the MAP satellite can measure the reionization optical depth $\\tau$ seventeen times more accurately using polarization information, and that polarization increases the sensitivity of the Planck satellite to tensor modes by a factor of 25.  Unfortunately, these large scales are also the ones where polarized foreground contamination is likely to be most severe, both because of the red power spectra of diffuse Galactic synchrotron and dust emission and because they require using a large fraction of the sky, including less clean patches. The key challenge in the CMB polarization endeavor will therefore be modeling, quantifying and removing large-scale polarized Galactic foregrounds. This is the topic of the present paper. We will use the POLAR experiment to provide the strongest limits to date on cross-polarized microwave background and foreground fluctuations on large angular scales, and employ polarization sensitive radio surveys to further quantify the polarized synchrotron radiation, which is likely to be the most challenging contaminant in the polarization maps expected from the MAP satellite.  \\begin{figure}[tb] \\preskip \\centerline{\\epsfxsize=9.0cm\\epsffile{polar2.ps}} \\postskip \\caption{\\label{summary}\\footnotesize% Summary of constraints on polarization so far. From top to bottom, the three curves show the concordance model predictions for $C_\\l^T$, $C_\\l^E$ and $C_\\l^X$, respectively. Four reionization models with $\\tau$=0.1, 0.2, 0.3 and 0.4 are also plotted (left dotted lines from bottom to top in both plots). The limits for $E$ are shown in the upper panel: Penzias \\& Wilson 65\\protect\\cite{PW65}, Caderni 78\\protect\\cite{C78}, Nanos   79\\protect\\cite{N79}, Lubin \\& Smoot 79\\protect\\cite{LS79}, Lubin \\& Smoot 81\\protect\\cite{LS81}, Sironi 98\\protect\\cite{S98}, Lubin  83\\protect\\cite{L83}, SASK (W93\\protect\\cite{W93},N97\\protect\\cite{N97}), TOCO (T99 hexagons\\protect\\cite{T99}), P88\\protect\\cite{P88}, F93\\protect\\cite{F93}, P97\\protect\\cite{P97}, S00\\protect\\cite{S00}, DMR\\protect\\cite{smootDMR}, PIQUE (H02\\protect\\cite{H02}) and POLAR (K01\\protect\\cite{keating01}). The limits for $X$ are shown in the lower panel: PIQUE (d0C02\\protect\\cite{angel_pique}) and POLAR (``This Work''). The shaded regions are the DASI results (Kv02\\protect\\cite{kovac02}). } \\end{figure}  At microwave frequencies, three physical mechanisms are known to cause foreground contamination: synchrotron, free-free and dust emission. When coming from extragalactic objects, this radiation is usually referred to as point source contamination and affects mainly small angular scales. When coming from the Milky Way, this diffuse Galactic emission fluctuates mainly on the large angular scales that are the focus of this paper. Except for free-free emission, all the above mechanisms are known to emit polarized radiation. In the near term, the best measurement of large-scale polarization will probably come from the MAP satellite. At MAP's frequency range (22-90 GHz), synchrotron radiation is likely to be the dominant polarized foreground \\cite{foregpars}. Unfortunately, we still know basically nothing about the polarized contribution of the Galactic synchrotron component at CMB frequencies \\cite{foregpars,tucci00,Baccigalupi00,Burigana02,bruscoli02,tucci02,giardino02}, since it has only been measured at lower frequencies and extrapolation is complicated by Faraday Rotation. This is in stark contrast to the CMB itself, where the expected polarized power spectra and their dependence on cosmological parameters has been computed from first principles to high accuracy \\cite{K97,ZS97,Z98,HuWhite97}.  Polarization of the Galactic continuum emission was first clearly detected in 1962 \\cite{Westerhout62}. In the succeeding years, polarization measurements of the northern sky were made at frequencies between 240 and 1415~MHz (see \\cite{S84} and references therein) with resolutions of only a few degrees. No large-area survey has been published since the compendium of Brouw and Spoelstra \\cite{BS76} and high-resolution surveys have only begun to be made recently. The first major investigation done after \\cite{BS76} is that of \\cite{Junkes87}, who observed a section of the Galactic plane defined by $49\\deg \\le \\ell \\le 76\\deg$ and $|b| \\le 15\\deg$, at frequency of 2.7~GHz. The study of \\cite{wieringa93} provides the highest resolution insight into the small-scale structure of the Galaxy; however, this only covered a few areas of the sky which were not larger than a degree or so across. Recently, two fully-sampled polarimetric surveys were done at 2.4~GHz \\cite{D95,D97} and 1.4~GHz \\cite{U98,U99}. All of these high-resolution surveys covered only regions near the Galactic plane, so in order to use them for inferences relevant to CMB experiments, they need to be extrapolated both in Galactic latitude and in frequency.  The rest of this paper is organized as follows. In section \\sec{phenomenology}, we review the basics of CMB and synchrotron polarization as well as our methods for measuring and modeling it. We present our results in \\sec{ResultsSec} and discuss our conclusions in \\sec{ConclusionsSec}.   \\begin{figure}[tb] \\preskip \\centerline{\\epsfxsize=7.5cm\\epsffile{ta_3fits.ps}} \\postskip \\postskip \\caption{\\label{tauFig}\\footnotesize% Examples of CMB polarization, showing how the reionization optical depth $\\tau$ affects the $T$ and $E$ power spectra (top) and the $TE$ correlation $r_\\l$ (bottom). Solid, dashed and dotted curves correspond to $\\tau$=0, 0.2 and 0.4, respectively. As discussed in \\protect\\cite{angel_pique}, changing the cosmological parameters affects the polarized and unpolarized power spectra rather similarly except for the cases of reionization and gravitational waves. In the reionization case, a new series of peaks are generated at large scales. $Top$-panel: Although there is no visible change in $T$ at large scales, there is clearly a visible change in $E$ since the Sachs-Wolfe nuisance is unpolarized and absent. $Lower$-panel: On small scales, reionization leaves the correlation $r_\\l$ unchanged since $C_\\l^T$ and $C_\\l^E$ are merely rescaled. On very large scales, $r_\\l$ drops since the new polarized signal is uncorrelated with the old unpolarized Sachs-Wolfe signal. On intermediate scales $\\l\\simgt 20$, oscillatory correlation behavior is revealed for the new peaks. For more details about CMB polarization and reionization see \\protect\\cite{Z97} } \\end{figure} ",
        "conclusions": "CMB polarization and its decomposition into $E$ and $B$ modes is a topic of growing importance and interest in cosmology. In the era of MAP, a key issue is to estimate the contribution of Galactic foregrounds (more specifically, polarized synchrotron emission) to these modes. We have used the POLAR experiment and radio surveys in order to quantify this contribution at large angular scales.  Using matrix-based quadratic estimator methods, we cross-correlated POLAR with DMR data and obtained upper limits of $E<7.4\\mu K$ and $|X|<11.1\\mu K$ % at 95\\% confidence. These upper limits are, unfortunately, to high to place intetesting constrains on reionization models. A similar cross-correlation analysis was performed by replacing the DMR with the Haslam data, obtaining an upper limit of $|X| \\simlt 15.4\\mu K$ % at 95\\% confidence.  We also used our quadratic estimator methods to measure the power spectra from the Leiden surveys, obtaining the following key results: \\begin{enumerate} \\item The synchrotron $E$- and $B$-contributions are equal to within 10\\% from 408 to 820~MHz, with a hint of $E$-domination at higher frequencies. One interpretation is that $E>B$ at CMB frequencies but that Faraday Rotation mixes the two at low frequencies. \\item Faraday Rotation \\& Depolarization effects depend not only on frequency but also on angular scale -- they are important at low frequencies ($\\nu\\simlt 10$ GHz) and on large angular scales. \\item We must take into account Faraday Rotation \\& Depolarization effects when extrapolating radio survey results from low to high galactic latitudes and from low to high frequencies. \\item We detect no significant synchrotron $TE$ cross correlation coefficient ($|r|\\simlt 0.2$), but Faraday Rotation could have hidden a substantial correlation detectable at CMB frequencies. \\item Combining the POLAR and radio frequency results, and the fact that the $E$-polarization of the abundant Haslam signal in the POLAR region is not detected at 30 GHz, suggests that the synchrotron polarization percentage at CMB frequencies is rather low. \\end{enumerate} Experiments such as MAP and Planck will shed significant new light on synchrotron polarization and allow better quantification of its impact both on these experiments and on ground-based CMB observations.  \\bigskip \\bigskip  This work was supported by NSF grants AST-0071213 \\& AST-0134999 and NASA grants NAG5-9194 \\& NAG5-11099. MT acknowledges a David and Lucile Packard Foundation fellowship and a Cottrell Scholarship from Research Corporation."
    },
    "0212/astro-ph0212075_arXiv.txt": {
        "abstract": "We present a new method for deriving cosmological constraints based on the evolution of the baryon mass function of galaxy clusters, and implement it using 17 distant clusters from our \\160d\\ \\emph{ROSAT} survey.  The method uses the cluster baryon mass as a proxy for the total mass, thereby avoiding the large uncertainties of the $M_{\\text{tot}}-T$ or $M_{\\text{tot}}-L_X$ relations used in all previous studies. Instead, we rely on a well-founded assumption that the $M_b/M_{\\text{tot}}$ ratio is a universal quantity, which should result in a much smaller systematic uncertainty. Taking advantage of direct and accurate \\emph{Chandra} measurements of the gas masses for distant clusters, we find strong evolution of the baryon mass function between $z>0.4$ and the present. The observed evolution defines a narrow band in the $\\Omega_m-\\Lambda$ plane, $\\Omega_m+0.23\\Lambda=0.41\\pm0.10$ at 68\\% confidence, which intersects with constraints from the Cosmic Microwave Background and supernovae~Ia near $\\Omega_m=0.3$ and $\\Lambda=0.7$. ",
        "introduction": "The growth of large scale structure in the near cosmological past is very sensitive to the value of the density parameter, $\\Omega_m$, and, weakly, to the cosmological constant, $\\Lambda$ (see, among others, Sunyaev 1971, Peebles 1980, Carroll, Press \\& Turner 1992). One of the strongest manifestations of the growth of density perturbations is the formation of virialized objects such as clusters of galaxies (Press \\& Schechter 1974). Therefore, the evolution of the cluster mass function provides a sensitive cosmological test (e.g.\\ Evrard 1989 and later works).  Present-day theory provides machinery for accurate predictions for the cluster mass function at any redshift, in any cosmology (see a brief review in \\S\\,\\ref{sec:theory}). Most of the difficulties with using cluster evolution as a cosmological probe are on the observational side.  The foremost problem is the necessity to relate cluster masses with other observables, such as the X-ray temperature. Even for the few brightest, best-studied clusters, the total mass within the virial radius --- the quantity generally used in theory --- is measured presently with large uncertainties (see, e.g., Markevitch \\& Vikhlinin 1997 and Fischer \\& Tyson 1997 for representative uncertainty estimates for the X-ray hydrostatic and weak lensing methods, respectively). To measure virial masses in a large, complete sample of clusters is presently impractical. Therefore, one usually resorts to other, more easily measurable quantities, which exhibit scaling relations with the total mass. The most widely used observable is the temperature of the intracluster medium. It is expected that $T$ is related to the virial mass as $M\\propto T^{3/2}$, and the available observations generally support this law (Nevalainen, Markevitch \\& Forman 2000; Horner, Mushotzky \\& Scharf 1999; Finoguenov, Reiprich \\& B\\\"ohringer 2001). To normalize this relation, virial masses still have to be measured in a representative sample of clusters. Any systematic uncertainties in the cluster mass measurements translate into the same uncertainty in the normalization of the $M-T$ relation.  The X-ray temperature functions (XTF) for both nearby and distant clusters have been used as a proxy for the mass function by a number of authors, using the normalization of the $M-T$ relation from different sources. An analysis of a statistically complete local XTF was first performed by Henry \\& Arnaud (1991). The analysis of the local XTF yields the normalization ($\\sigma_8$), in combination with $\\Omega_m$, and slope of the power spectrum of density fluctuations on cluster scales. This type of study has since been applied to ever improving observational data (see Evrard et al.\\ 2002 for a recent summary).  A first cosmological measurement using the evolution of the XTF at redshifts greater than zero is due to Henry (1997), who derived $\\Omega_m \\approx 0.5\\pm0.15$. A similar analysis was performed by Eke et al.\\ (1998), Donahue \\& Voit (1999), Henry (2000), Blanchard et al.\\ (2000), with updates on both observational and theoretical sides as well as with different normalizations of the $M-T$ relation.  We also note the studies by Reichart et al.\\ (1999) and Borgani et al.\\ (2001) who modeled the X-ray luminosity functions for distant clusters using an empirical $L-T$ relation and the usual theory for the XTF evolution.  Theoretical prediction for the number density of clusters with a given temperature is exponentially sensitive to the normalization of the $M-T$ relation.  A 30\\% change in the normalization of the $M-T$ relation results in a 20\\% change in the determination of the $\\sigma_8$ parameter and shifts an estimate of $\\Omega_m$ by a factor of 1.3--1.4 (e.g., Rosati, Borgani \\& Norman 2002). Given that $\\pm50\\%$ is a fair estimate of the present systematic uncertainties in the $M-T$ normalization, the impact of this relation on cosmological parameter determination can be very significant.  In modeling the high-redshift temperature function, the theoretically expected evolution of the $M-T$ relation is usually assumed: $M\\propto T^{3/2}/(1+z)^{3/2}$. It would be very important to verify this evolution at high $z$. Unfortunately, this seems to push the limits of current observational techniques. The evolution of other scaling relations can be also quite important. For example, the $L-T$ relation is required for computing the volume of X-ray flux limited surveys, and for estimating cluster masses, when the X-ray luminosity function is used as a proxy for the mass function. Most of the previous studies have assumed that the $L-T$ relation does not evolve, however recent \\emph{Chandra} measurements indicate otherwise (Vikhlinin et al.\\ 2002).  Finally, most of the available cosmological studies are based on the distant cluster sample from the \\emph{Einstein} Extended Medium Sensitivity Survey (Gioia et al.\\ 1990, Henry et al.\\ 1992). The completeness of this survey at high redshifts is often questioned (Nichol et al.\\ 1997; Ebeling et al.\\ 2000; Lewis et al.\\ 2002). Although we believe that the EMSS is essentially complete (as confirmed by good agreement between the XLFs from the EMSS and several \\emph{ROSAT} surveys --- Gioia et al.\\ 2001, Vikhlinin et al.\\ 2000, Mullis et al., in preparation), it is arguably important to try another completely independent sample for the cosmological studies based on cluster evolution.  The purpose of this paper is to apply the evolutionary test to a new sample of distant, $z>0.4$, clusters derived from the 160~deg$^{2}$ \\emph{ROSAT} survey (Vikhlinin et al.\\ 1998), relying on the evolution of the cluster scaling relations as measured by \\emph{Chandra}, and using a modeling technique which does not rely upon the total mass measurements to normalize the $M-T$ or $M-L$ relations and therefore bypasses many of the uncertainties mentioned above.  The proposed method relies on the assumption that the baryon fraction within the virial radius in clusters should be close to the average value in the Universe, $f_{b,U} = \\Omega_b/\\Omega_m$. This is expected on general theoretical grounds (White et al.\\ 1993) because gravity is the dominant force on cluster scales. The universality of the baryon fraction in massive clusters is supported by cosmological numerical simulations (e.g.\\ Bialek, Evrard \\& Mohr 2001) and by most available observations (most recently, Allen et al.\\ 2002). In principle, the equality $f_b = \\Omega_b/\\Omega_m$ can be used to determine cosmological parameters either from the absolute measurements of $f_b$ in the nearby clusters (White et al.\\ 1993) or from the apparent redshift dependence of $f_b$ measurements in the high-redshift clusters (Sasaki 1996, Pen 1997). Both methods, however, require an accurate observational determination of the total mass in individual clusters, which we would like to avoid.  Unlike the total mass, the baryon mass in clusters is relatively easily measured to the virial radius (Vikhlinin et al.\\ 1999). To first order, the baryon and total mass are trivially related, $M_b = M\\,\\Omega_b/\\Omega_m.$ If this is the case, the relation between the cumulative total mass function, $F(M)$, and the baryon mass function, $F_b(M_b)$, is also trivial, \\begin{equation}\\label{eq:f(m_b):f(m):ideas} F_b(M_b) = F\\left(\\Omega_m/\\Omega_b\\, M_b\\right). \\end{equation} The average baryon density in the Universe is accurately given by observations of the light element abundances and Big Bang Nucleosynthesis theory (Burles, Nollet \\& Turner 2001). Therefore, we have everything needed to convert the theoretical model for the total mass function, $F(M)$, to the prediction for a directly measurable baryon mass function --- to compute $F(M)$ one must choose a value of $\\Omega_m$, which fixes the scaling between the total and baryon mass. The conversion of the total to the baryon mass function --- at least to a first approximation --- can be expressed through the model parameters, and therefore $F_b(M_b)$ can be modeled with the usual set of parameters, $\\Omega_m$, $\\sigma_8$, $n$, $h$, and at a high redshift --- $\\Lambda$. We will show below that the evolution of the baryon mass function provides robust constraints on a combination of $\\Omega_m$ and $\\Lambda$. Modeling of the baryon mass function at $z=0$ provides a measurement of the power spectrum normalization, $\\sigma_8$, and of the shape parameter $\\Gamma=\\Omega_m\\,h$, with degeneracies that differ from those given by the temperature function (Voevodkin \\& Vikhlinin, in preparation).  Theoretical models of the cluster mass function usually consider masses measured within radii defined by certain values of the density contrast, $\\delta = 3 M(r)/(4\\pi\\,r^3)/\\rho$, where $\\rho$ is either the critical density or mean density of the Universe. The assumption that the baryon fraction in clusters is universal at sufficiently large radii allows one to determine the overdensity radii $r_\\delta$ without measuring the total mass. Indeed, the baryon and matter overdensities are then equal, $\\delta_b=\\delta$, and the baryon overdensity is defined by a baryon mass $M_b(r)$ and the mean baryon density in the Universe, known from the BBN. One of the most accurate theoretical models for the mass function (Jenkins et al.\\ 2001) uses masses corresponding to the overdensity $\\delta=324$ relative to the mean density of the Universe at the redshift of observation. The matching baryon masses, $M_{b,324}$ are easily measured using \\emph{ROSAT} data for nearby clusters (Voevodkin, Vikhlinin \\& Pavlinsky 2002a, VVPa hereafter), and \\emph{Chandra} data for distant clusters (Vikhlinin et al.\\ 2002).  The paper is organized as follows. In \\S\\,\\ref{sec:data}, we discuss the distant cluster sample and the gas mass measurements in the distant and nearby clusters. The baryon mass function for distant clusters is derived in \\S\\,\\ref{sec:mfun:meas}. The theory for modeling these data is reviewed in \\S\\,\\ref{sec:theory}. We describe our fitting procedure in \\S\\,\\ref{sec:fitting}, and present the derived constraints on $\\Omega_m$ and $\\Lambda$ in \\S\\,\\ref{sec:omega:lambda:results}.  All cluster parameters are quoted for $h=0.65$, $\\Omega_m=0.3$, and $\\Lambda=0.7$. The X-ray fluxes and luminosities are in the 0.5--2 keV energy band. Measurement uncertainties are $1\\sigma$. ",
        "conclusions": "We have presented a new method of applying the models of the cluster evolution to observations at high redshifts. This method is based on easily observable baryon mass function. Unlike the studies based on the cluster temperature or luminosity functions, our method does not suffer from large uncertainties in the $M_{\\text{tot}}-T$ and $M_{\\text{tot}}-L$ relations. It uses only a physically motivated assumption that the baryon fraction in the cluster mass within the large radii is close to the universal value, $\\Omega_b/\\Omega_m$.  Systematic uncertainties in this method are smaller and very different from those using the temperature or luminosity functions. Until improved observational techniques are able to normalize the scaling relations to a $10\\%$ accuracy at all redshifts, modeling of the baryon mass function appears to be a promising method for deriving the cluster-based cosmological constraints.  The application of this method to the subsample of distant clusters from the \\160d\\ survey results in a tight constraint on the combination of cosmological parameters $\\Omega_m$ and $\\Lambda$. This constraint intersects the current CMB and SN~Ia constraints in a small area near $\\Omega_m=0.3$ and $\\Lambda=0.7$, showing a remarkable agreement of totally independent methods.  We note an excellent agreement of our cluster-derived cosmological parameter $\\Omega_m$ with the results based on the X-ray measurements of the total mass and baryon fraction in the intermediate-redshift clusters (Allen et al.\\ 2002). Also, there is a good agreement with those studies of the XTF evolution in the EMSS sample which use the normalization of the $M-T$ relation based on the X-ray mass measurements (e.g., Henry 2000). Such consistency indirectly confirms the validity of the X-ray cluster mass measurements.  Our method relies on the assumption that the baryon fraction in massive clusters does not evolve with redshift. This assumption is well-supported by numerical simulations. One issue of concern, however, is that these same simulations do not reproduce the observed evolution of the $M_b-T$ relation. At a fixed $T$, we observe $M_b \\propto (1+z)^{-0.5\\pm0.4}$ for $\\Omega_m=0.3$ and $\\Lambda=0.7$ (Vikhlinin et al.\\ 2002), which is significantly different from the expected evolution $(M_b, M_{\\text{tot}})\\propto T^{3/2}(1+z)^{-3/2}$ (e.g., Bryan \\& Norman 1998). This may indicate that the baryon fraction in clusters evolves with redshift. However, it is much more likely that simulations simply do not reproduce the distribution of gas in the cluster central regions. These regions usually dominate in the observed emission-weighted $T$ but at the same time the gas there is especially prone to processes such as radiative cooling which can easily change $T$ but are unlikely to modify $M_b$ or $M_{\\text{tot}}$.  The accuracy of our cosmological constraints is presently limited by the low number of the distant clusters. When new distant cluster samples from the \\emph{Chandra}, \\emph{XMM}, and extended \\emph{ROSAT} surveys become available, and more distant clusters are observed with adequate exposures by \\emph{Chandra} and \\emph{XMM}, even tighter cosmological constraints will be possible. Eventually, different methods (cluster evolution, CMB, SN~Ia etc.) will start to disagree, which will lead to better understanding of the systematics and warrant more complicated theoretical models, e.g.\\ those involving non-standard equations of state for the dark energy."
    },
    "0212/astro-ph0212243_arXiv.txt": {
        "abstract": "Collisions and mergers of gas-rich galaxies trigger bursts of star and cluster formation.  Of the thousands of clusters typically formed during a major merger, only the most massive and compact survive for Gigayears as globular clusters (GCs).  In $\\la$1~Gyr old merger remnants, these `second-generation' GCs appear by the hundreds as young halo clusters of $\\sim$\\,solar metallicity.  Their likely descendants, metal-rich GCs of intermediate age (2\\,--\\,5~Gyr), have recently been found in $\\sim$10 E galaxies, where they appear as slightly overluminous red GCs with a still power-law-like luminosity function.  Their color and radial distributions suggest that they evolve into the red metal-rich GCs observed in old ellipticals.  There is good evidence that second-generation GCs form from giant molecular clouds shocked by the rapid pressure increase in merger-induced starbursts.  This mechanism supports the view that the universal pressure increase during cosmological reionization may have triggered the formation of the metal-poor globulars observed in galaxies of all types. ",
        "introduction": "Ivan King was my revered teacher.  I owe him a lot, including an interest in globular clusters. Into my copy of his textbook {\\it The Universe Unfolding} (1976), he wrote the dedication:  ``Make as much of this obsolete as you can, Fran\\c cois! Best wishes, Ivan.''  In this spirit I put forth the hypothesis that major mergers of gas-rich spirals form not only ellipticals, but also new globular clusters within them (Schweizer 1987).  I cherish Ivan's later comment on this hypothesis: ``When Schweizer suggested resolving the $S_N$ problem by making globular clusters in mergers, I thought it was ridiculous.  But now in the HST observations of NGC 1275 we seem to see it actually happening'' (King 1993).  Always quick to grasp major issues, Ivan challenged Alar Toomre at the Yale Conference as follows (King 1977): ``You showed us 10 merging pairs [of galaxies] and then asked us to look for, or at least accept the existence of, 500 remnants from so long ago that they no longer bear the `made by Toomre' label.  I would be much more impressed if you showed us the 20 or 30 systems in the box immediately adjacent in your histogram.  What do these merged pairs look like in their next few galactic years?''  As I hope to show in the present review, the descendants of major disk--disk mergers---i.e., ellipticals with globular clusters of intermediate age---are now being found in growing numbers. ",
        "conclusions": ""
    },
    "0212/astro-ph0212525_arXiv.txt": {
        "abstract": "The second largest \\ion{H}{2} region in the Large Magellanic Cloud, N11B has been surveyed in the near IR. We present $JHKs$ images of the N11B nebula. These images are combined with CO$(1\\rightarrow0)$ emission line data and with archival NTT and HST/WFPC2 optical images to address the star formation activity of the region. IR photometry of all the IR sources detected is given. We confirm that a second generation of stars is currently forming in the N11B region. Our IR images show the presence of several bright IR sources which appear located towards the molecular cloud as seen from the CO emission in the area. Several of these sources show IR colours with YSO characteristics and they are prime candidates to be intermediate-mass Herbig Ae/Be stars. For the first time an extragalactic methanol maser is directly associated with IR sources embedded in a molecular core. Two IR sources are found at $2''$ (0.5 pc) of the methanol maser reported position. Additionally, we present the association of the N11A compact \\ion{H}{2} region to the molecular gas where we find that the young massive O stars have eroded a cavity in the parental molecular cloud, typical of a champagne flow. The N11 region turns out to be a very good laboratory for studying the interaction of winds, UV radiation and molecular gas. Several photodissociation regions are found. ",
        "introduction": "Giant \\ion{H}{2} regions are relatively scarce objects in the Local Group. In these regions we expect to find a broad spectrum of coeval phenomena related to massive stars. The action of stellar winds from the massive stars and the supernova explosions pushes and destroys the natal molecular cloud, but also, helps to trigger the formation of a new generation of massive stars. At present, there is morphological evidence showing that a former massive star generation could produce a new one in the peripheral molecular clouds, but there is no quantitative evidence of how different generations are related in a giant \\ion{H}{2} region. This quantitative evidence is very hard to obtain, because it requires the characterisation of the stellar content embedded in the \\ion{H}{2} region and the physical condition of the gas in the molecular cloud.  The sequential star formation scenario proposed by Elmegreen \\& Lada (1977, see also Elmegreen 1998), has a few clear galactic examples where it is possible to see that the direct action of massive stars in the parental molecular cloud could be producing a new generation of stars. Two such regions, the OB Association Ara OB1 (Arnal et al. 1987), and the Carina Nebula (Smith et al. 2000) have been proposed as leading cases in our Galaxy.  In general, in a massive star forming region we find a cavety of relatively low density ionised gas produced by the action of several hot massive stars together with photodissociation regions (PDR) located on the surfaces of the molecular clouds facing the hot stars. Embedded IR sources and/or molecular cores could be direct evidence of an emerging  new generation of stars.  In the Local Group, the formation and evolution of large superbubbles (Oey 1999) is a controversial issue. Evolved superbubbles are suggested to be produced by the direct action of SN explosions in the interstellar medium which produce shells of emitting gas (Wang \\& Helfand 1991). In younger \\ion{H}{2} superbubbles, the main bright nebular filaments are not shells but PDR interfaces between the \\ion{H}{2} cavity and the surrounding molecular clouds surfaces. An example in an extragalactic environment is now clearly established: 30 Doradus Nebula in the Large Magellanic Cloud (Rubio et al. 1998; Walborn et al. 1999a and references therein).  The N11 nebular complex (Henize 1956) or DEM34 (Davies, Elliot, \\& Meaburn 1976) is the second largest \\ion{H}{2} region in the LMC after 30 Doradus Nebula (Kennicutt \\& Hodge 1986), lying at the opposite end of the LMC Bar. The complex consists of a huge bubble surrounded by nine distinct nebular entities (Rosado et al. 1996). These are easily distinguished in the spectacular H$\\alpha$+[\\ion{N}{2}] emission line image published by Walborn \\& Parker (1992). The OB association LH9 is located at the center of a cavity of about $80\\times60$ pc (Lucke \\& Hodge 1970; Parker et al. 1992). This OB association is dominated by HD\\,32228 (Radcliffe\\,64, Sk\\,$-66\\,28$, Breysacher\\,9), a massive compact cluster containing a WC4 star of about 3.5 Myr (Walborn et al. 1999b). Surrounding LH9 are several younger OB associations embedded in dense nebular regions, two of which have O3 stars among their members: LH13 in N11C (Heydari-Malayeri et al. 2000) and LH10 in N11B (Parker et al. 1992). The latter authors propose an evolutionary link between the OB associations LH9 and LH10, where the star formation in LH10 could have been triggered by the evolution of massive stars in LH9. This suggestion was based on the different initial mass function slopes found, significantly flatter for LH10 than for LH9, and the different apparent ages derived for both associations. Walborn \\& Parker (1992) noticed a remarkable analogy in the structural morphology between N11 and 30 Doradus and proposed a two-stage starburst scenario to explain the morphological distribution of OB stars in both \\ion{H}{2} regions. According to this scenario, an initial and centrally concentrated burst of stars may trigger a second burst in the peripherical molecular clouds about $2\\times10^6$ years later. In N11, this process would be 2 million years older than in 30 Dor.  N11 is a large molecular cloud complex composed by at least 29 separate molecular clouds (Israel et al. 2002, Israel \\& de Graauw 1991), directly associated with the ionised gas (Rosado et al. 1996). The CO emission is concentrated in distinctive peaks and correlated with the H$\\alpha$ and FIR dust emissions, suggesting that the star-formation activity is distributed in the ringlike structure of N11 (Caldwell \\& Kutner 1996).  Thus, the N11 region is a very good candidate for the study of the sequential star formation processes and for the determination of how different complex elements are related in the ongoing star generation embedded in the nebular ring.  In this paper we present new near-infrared images and CO observations of the N11B nebula which show a striking relationship between IR sources and the molecular cloud. In addition, we compare our observations with optical nebular emission archival images obtained with the New Technology Telescope (NTT) and the Hubble Space Telescope (HST) yielding further evidence of  the ongoing star forming activity in this region. ",
        "conclusions": "We confirm that a second generation of stars is currently forming in the N11B region. Our IR images show the presence of several bright IR sources which appear located toward the molecular cloud as seen from the CO emission in the area. Several of  these sources show IR colours with YSO characteristics and they are prime candidates to be intermediate-mass Herbig Ae/Be stars following the  criterion of Brandner et al. (2001). For the first time, an extragalactic methanol maser is directly associated with IR sources embedded in a molecular core. Two IR sources, BRRG\\,147 and 148, are found at $2''$ (0.5 pc) of the methanol maser reported position. They are among the reddest IR sources in N11B and therefore the strongest candidates of YSOs. IR spectroscopy of these sources will reveal their nature.  The overall picture of N11B including the CO and ionised gas distributions clearly demonstrate that the nebula is a region of interaction between stellar winds and UV radiation from hot young stars, as well as of PDRs in the molecular clouds originated by the action of such stars. Almost all the optical nebular emission in N11B is directly associated with CO emission, suggesting that the main source of such optical emission comes from the molecular cloud surface, where the gas is photoevaporated and ionised. In fact, all the brighter optical nebular emission filaments and peaks are PDRs facing the hot stars: a bright rim $10''$ to the north of the multiple O star PGMW\\,3120, a Y-shaped cometary globule close to the double O star PGMW\\,3223, a kiwi-shaped globule $10''$ to the west of the multiple early-O star PGMW\\,3204/09  The nebular knot N11A to the east of N11B is also associated with a molecular cloud core. This knot is excited by the multiple O star PGMW\\,3264, and we find evidence of the interaction of the stellar winds and the molecular gas. This nebula has the morphology and kinematics of a champagne flow where the massive stars have open a cavity in their parental molecular cloud. Future IR data will allow to search for IR embedded sources in the region.  Walborn \\& Parker (1992) proposed that N11 was few $10^5$ years older than the 30 Doradus, based on their analysis of the relative ages of the stellar content of both giant \\ion{H}{2} regions. In 30 Doradus, a new generation of massive stars is currently being formed in the peripherical molecular clouds apparently due to the energetic activity of the hot massive stellar core R136 (Rubio et al. 1998; Walborn et al. 1999a; Brandner et al. 2001). The stellar content of such new generation in 30 Doradus has O stars embedded in compact nebulosities (Walborn et al. 2002 and references therein), with a relative age of about 1-2 Myr younger than the core R136, characterised by the presence of luminous H-burning WN stars with an age of about 2 Myr. Walborn et al. (1999b) derived an age of about 3.5 Myr for the LH9 association in the N11 core, and less than 1 Myr for PGMW\\,3209 in LH10. Their results confirm the youth of the LH10 association (also deduced from absence of WR stars), and that the N11 region is older than 30 Doradus.  Our study supports the idea that N11B is in a later stage of evolution than 30 Doradus. It shows a lack of bright IR sources compared to the star forming regions in 30 Dor. Our brightest IR source has $K=15$, while several $K=12$ IR embedded stars were observed in 30 Dor. The IR sources in N11B are probably intermediate mass star candidates still embedded in the molecular gas. The O stars have blown away the molecular material and are disrupting the molecular cloud surface. Thus, we may be witnessing the last stage of the second generation burst in N11B."
    },
    "0212/astro-ph0212239_arXiv.txt": {
        "abstract": "We have analyzed ISO SWS\\,01 observations for 61 proto-planetary ~nebulae ~candidates ~and ~classified ~their spectra according to their dominant chemistry. On the basis of our classification and the more general classification of SWS\\,01 spectra by Kraemer et al.\\,(2002) we discuss the connection between  proto-planetary nebulae candidates and planetary nebulae, with emphasis on possible precursors of planetary nebulae with [WR] central stars. ",
        "introduction": "The proto-planetary nebula (PPN) phase, in the context of the late stages of low- and intermediate-mass star evolution, is a short lasting period (of the order of at most a few thousands of years) when stars evolve from Asymptotic Giant Branch (AGB) to planetary nebula (PN). The PPN phase is characterized by a rather small mass loss rate ($\\sim$10$^{-7}$\\,M$_{\\odot}$yr$^{-1}$), a decrease of the circumstellar shell optical depth due to expansion, and a decrease of the star radius and consequent increase of the effective temperature (typically from about 4000\\,K until the onset of ionization at about 25000\\,K) due to gradual consumption of stellar envelope, both by nuclear processing and by stellar wind.  PPNe, being immediate precursors of PNe deserve special attention because they offer the possibility to identify the main physical and chemical processes which lead to diversity of shapes and chemical compositions in PNe and their central stars. One of the most intriguing class of PNe is the class of planetary nebulae with Wolf-Rayet type central stars (hereafter [WR]PNe -- see e.g. Tylenda et al.\\,1993, G\\'orny \\& Stasi\\'nska\\,1995, Tylenda\\,1996, Leuenhagen \\& Hamann\\,1998). Among about 1500 galactic PNe $\\sim$50 are [WR]PNe (G\\'orny \\& Tylenda\\,2000). Observations with the Infrared Space Observatory (ISO, Kessler et al. 1996) have shown that in [WR]PNe both forms of dust (C-rich: PAHs and O-rich: crystalline silicates) are present (Waters et al.\\,1998a, Cohen et al.\\,1999). For a discussion of scenarios put forward to explain this unexpected discovery see recent papers by Cohen et al. (2002), De Marco et al.\\,(2002) and De Marco \\& Soker\\,(2002).  In this paper we report on a search of the ISO Data Archive (IDA) for spectra of PPNe taken with the Short Wavelength Spectrometer (SWS\\,01, de Graauw et al. 1996) and present a classification of these spectra. We then briefly discuss the relation between PPNe and PNe based on the results of this classification and the one performed by Kraemer et al.\\,(2002). ",
        "conclusions": "\\begin{table}[tb]{} \\caption[]{Source classification} \\begin{center} \\leavevmode \\footnotesize \\begin{tabular}{l @{   } l @{  } l | l @{  } l @{   } l }\\hline IRAS         & KSPW$^1$      & {\\it c21} & IRAS         & KSPW      & {\\it c21} \\\\ name         & class     &           & name         & class     &  \\\\ \\hline \\noalign{\\smallskip} \\multicolumn{3}{}{}C\\,21 \\\\ \\noalign{\\smallskip} \\object{16594$-$4656}  & 4.CT & 7.72 & \\object{20000$+$3239}  & 4.CT  & 3.68 \\\\ \\object{23304$+$6147}  & 4.CT & 7.13 & \\object{Z02229$+$6208} & 4.CT  & 2.96 \\\\ \\object{19500$-$1709}  & 4.CT & 5.18 & \\object{22272$+$5435}  & 4.CT  & 2.94 \\\\ \\object{07134$+$1005}  & 4.CT & 5.14 & \\object{05341$+$0852}  & 4.PN  & 1.79 \\\\ \\object{22574$+$6609}  & 4.PUp:    & 4.09 &               &           &      \\\\ \\noalign{\\smallskip} \\multicolumn{3}{}{}C\\,mol \\\\ \\noalign{\\smallskip} \\object{RAFGL 2688}$^2$& 4.CN & 7.15 & \\object{01144$+$6658} & 4.CR & 1.79 \\\\ \\object{23321$+$6545} & 4.CN: & 4.69 & \\object{19548$+$3035} & 4.CR & 1.57 \\\\ \\object{19480$+$2504} & 4.F   & 3.17 &              &           &      \\\\ \\noalign{\\smallskip} \\multicolumn{3}{}{}C\\,RCrB \\\\ \\noalign{\\smallskip} \\object{15465$+$2818} & 3.W  & 0.45 & \\object{14316$-$3920} & 3.W   & 0.28 \\\\ \\object{19132$-$3336} & 3.W  & 0.29 &              &           &      \\\\ \\noalign{\\smallskip} \\multicolumn{3}{}{}C\\,PAH \\\\ \\noalign{\\smallskip} \\object{13428$-$6232} & 4.Fu   & 15.1 & \\object{13416$-$6243} & 4.CN:u & 3.34 \\\\ \\object{01005$+$7910} & 4.PUp  & 4.66 & \\object{06176$-$1036} & 4.U/SC & 1.24 \\\\ \\object{17347$-$3139} & 4.UE   & 4.62 & \\object{16235$-$4832} & 3.CR:: & 0.74 \\\\ \\object{16279$-$4757} & 4.U/SC & 3.82 & \\object{10158$-$2844} & 2.U    & 0.21 \\\\ \\noalign{\\smallskip} \\multicolumn{3}{}{}Si\\,A \\\\ \\noalign{\\smallskip} \\object{15553$-$5230} & 4.SA: & 7.01 & \\object{15452$-$5459} & 4.SAe & 3.55 \\\\ \\object{18276$-$1431} & 4.SA  & 6.45 & \\object{17195$-$2710} & 4.SA  & 2.97 \\\\ \\object{17150$-$3224} & 4.SA  & 6.43 & \\object{19386$+$0155} & 4.SB  & 2.57 \\\\ \\object{22036$+$5306} & 5.SA  & 5.49 & \\object{19343$+$2926} & 5.SA  & 2.44 \\\\ \\object{18596$+$0315} & 5.SA: & 3.84 & \\object{17516$-$2525} & 4.SAp & 1.89 \\\\ \\noalign{\\smallskip} \\multicolumn{3}{}{}Si\\,E \\\\ \\noalign{\\smallskip} \\object{12175$-$5338} & 4.SE:  & 7.98 & \\object{11385$-$5517} & 5.SE: & 1.49 \\\\ \\object{18095$+$2704} & 4.SEC  & 2.79 & \\object{20004$+$2955} & 2.SEc & 3.55 \\\\ \\object{18062$+$2410} & 4.SE:  & 2.59 & \\object{17534$+$2603} & 3.SEp & 0.46 \\\\ \\object{19244$+$1115} & 4.SEC  & 1.71 &              &            &      \\\\ \\noalign{\\smallskip} \\multicolumn{3}{}{}Si\\,RVTau \\\\ \\noalign{\\smallskip} \\object{18281$+$2149} & 4.SEC  & 1.49 & \\object{18448$-$0545} & 2.SEa: & 0.45 \\\\ \\object{22327$-$1731} & 4.SE:: & 1.30 & \\object{12185$-$4856} & 7  R   & 0.35 \\\\ \\object{20117$+$1634} & 4.SE:: & 0.56 &              &            &     \\\\ \\noalign{\\smallskip} \\hline \\multicolumn{2}{}{} $^1$--Kraemer et al.\\,(2002);\\\\ \\multicolumn{5}{}{} $^2$--RAFGL\\,2688 has not been observed by IRAS.\\\\ \\end{tabular} \\end{center} \\end{table}  From Table\\,1, the proportion of C-rich  sources is 25/47 (53\\,\\%). However, because of source selection, the ISO sample is biased towards sources with the 21\\,$\\mu$m feature, which amount to 9 objects in our sample. Therefore, the unbiased proportion of C-rich PPNe candidates is somewhere between 16/38 (42\\,\\%) and  53\\,\\%. Interestingly, the proportion of C-rich PNe is 35--40\\,\\% (Rola \\& Stasi\\'nska\\,1994, Kingsburgh \\& Barlow\\,1994). This suggests that most PPNe candidates from our sample will indeed become PNe (from now on we will drop the word ``candidate'' for simplicity).  We can now roughly estimate what number of PPNe from our classified ISO sample are expected to become [WR]PNe.  G{\\'o}rny \\& Stasi{\\'n}ska\\,(1995) argued that the proportion of [WR]PNe relative to the total number of PNe is about 8\\,\\%. Then one expects the same proportion of [WR]PNe precursors among the sample of PPNe. This implies that 8\\,\\% of 47 PPNe, i.e $\\sim$4 will become [WR]PNe.  As discussed recently by De Marco \\& Soker\\,(2002) the most important characteristic of [WR]PNe seems to be dual dust chemistry.  About 80\\,\\% of {\\em late} [WR]PNe show the presence of silicates -- mostly crystalline -- while at the same time all of them show PAH emission, see Szczerba et al.\\,(2001a). Very likely the dual dust chemistry is a result of O-rich dust being formed when the star was O-rich and stored in some stable reservoir in the stellar vicinity, while the carbon chemistry occurs later when the [WR]PN progenitor is already C-rich. As concerns {\\em early} [WR]PNe, crystalline silicates are seen in only one object (NGC\\,5315 -- K. Volk, private communication) out of 6 with available ISO spectra. The reason for such a small proportion is not clear. G{\\'o}rny \\& Tylenda\\,(2000) argued that there is an evolutionary link between late and early [WR]PNe, so a stable  reservoir of crystalline silicates should be seen in later phases as well, unless a fast wind is able to destroy or remove the crystalline material from the star surroundings. Non-detection of crystalline silicates can also be  due to a worse quality of ISO SWS data for early [WR]PNe. From the above considerations, it is natural to think that PPNe showing dual dust chemistry should evolve into [WR]PNe. In our sample, the famous Red Rectangle (IRAS\\,06176$-$1006) and  IRAS\\,16279$-$4757  contain both C-rich dust and crystalline silicates (from the KSPW classification, see also Waters et al.\\,1998b and Molster et al.\\,1999) which makes them good candidates to be precursors of late [WR]PNe. Note that not all [WR]PNe are C-rich (G\\'orny \\& Stasi\\'nska 1995, De Marco\\,2002), in consequence not all PPNe that will become [WR]PNe are necessarily C-rich. Therefore, further candidates for [WR] PNe precursors in our sample are then the O-rich sources AC\\,Her (IRAS\\,18281$+$2149), IRAS\\,18095$+$2704 and IRAS 19244$+$1115.  If non-detection of dual dust chemistry in some [WR] PNe is not due to observational effects but is a sign of real absence of crystalline silicates, then it is possible that the precursors of these [WR]PNe may also not show crystalline silicate features. Recently, Hony et al.\\,(2001) reported that some [WR]PNe show the 21\\,$\\mu$m feature suggesting a possible link between them and PPNe with the 21\\,$\\mu$m feature (our class C\\,21 or KSPW class CT). To our knowledge none of the C\\,21 source shows evidence of dual dust chemistry so they could be candidates for precursors of [WR]PNe  without crystalline silicates. Demographic arguments show that this could be the case for only a small fraction of them. Indeed, there are at least 12 PPNe with the 21\\,$\\mu$m feature (see e.g. Kwok et al.\\,2002) among 220 known PPNe, i.e. 5.5\\%. On the other hand only 20\\,\\% of late type [WR]PNe with analyzed ISO spectra do not have crystalline silicates. According to G\\'{o}rny \\& Tylenda\\,(2000), 30\\,\\% of all [WR]PNe are of late type and, as mentioned above, [WR]PNe represent 8\\% of the total number of PNe. Thus, the proportion of {\\em late} [WR]PNe without crystalline silicates is {\\em at most} 20\\%$ \\times$ 30\\%$ \\times$ 8\\% = 0.5\\% of the total population of PNe. Therefore, most  of 21\\,$\\mu$m sources will not go through the  {\\em late} [WR]PN phase. One could still argue that they could evolve into PNe in which the [WR] phenomenon will appear only later. However, G\\'{o}rny \\& Tylenda (2000) have shown that most [WR]PNe do evolve from late to early type. Besides, among PNe, the 21\\,$\\mu$m feature has been seen only in [WR]PNe, meaning that PPNe with the 21\\,$\\mu$m feature cannot evolve to early [WR]PNe that have not gone through a late [WR] stage. Therefore, PPNe with the  21\\,$\\mu$m feature cannot be considered, as a class, as precursors of [WR]PNe."
    },
    "0212/astro-ph0212149_arXiv.txt": {
        "abstract": "s{ Abundance  observations\tindicate the presence of rapid-neutron capture (i.e., {\\it r}-process) elements in old Galactic halo and globular cluster stars. These observations demonstrate that the earliest generations of stars in the Galaxy, responsible for neutron-capture synthesis and the progenitors of the halo stars, were rapidly evolving. Abundance comparisons among large numbers of stars provide clues about the nature of neutron-capture  element synthesis both during the earliest times and throughout the history of the Galaxy. In particular, these comparisons suggest differences in the way the heavier (including Ba and above) and lighter neutron capture elements are synthesized in nature. Understanding these differences will help to identify the astrophysical site (or sites) of and conditions in the {\\it r}-process. The abundance comparisons also demonstrate a large star-to-star scatter in the neutron-capture/iron ratios at low metallicities- which disappears with increasing [Fe/H]- and suggests an early, chemically unmixed and inhomogeneous Galaxy. The  very recent neutron-capture element observations indicate  that the early phases of Galactic nucleosynthesis, and the associated chemical evolution,  are  quite complex, with the yields from different (progenitor) mass-range stars contributing to  different chemical mixes. Stellar abundance comparisons  suggest  a change from the {\\it r}-process to the slow neutron capture ({\\it i.e.}, {\\it s}-) process at  higher metallicities (and later times)in the Galaxy. Finally, the detection of thorium and uranium in halo and globular cluster stars offers a promising, independent age-dating technique that can put lower limits on the age of the Galaxy and thus the Universe.} ",
        "introduction": "The heavy solar system abundances (here, Z $>$ 30) are formed in neutron capture ({\\it n}-capture) processes, either the slow ({\\it s}-) or rapid ({\\it r}-) process. We show in Figure~\\ref{solar} the solar system -- also thought of as cosmic -- abundances with the neutron-capture elements highlighted. Abundance observations of these elements in halo stars contain vital clues  to the nucleosynthesis history and chemical evolution of the Galaxy, and abundances  of radioactive elements can also be utilized to obtain age determinations  for the oldest stars, which in turn put lower limits on age estimates for the Galaxy and the Universe ({\\it e.g.}, [1,2]).  \\begin{figure}[ht] \\centerline{\\epsfxsize=3.9in\\epsfbox{solar.eps}} \\caption{ Abundances of elements in the Sun and in undifferentiated solar system material [3]. This abundance set is normalized by convention to log N(H) = 12 in astronomical literature on stars (planetary and meteoritic studies usually normalize the abundances to log~N(Si)~= 6). The main figure shows the entire set of stable and long-lived radioactive elements, while the inset is restricted to the $n$-capture elements, defined here as those elements with Z~$>$ 30. \\label{solar}} \\end{figure} ",
        "conclusions": ""
    },
    "0212/astro-ph0212180_arXiv.txt": {
        "abstract": "I describe a unique, 20-year-long timing program for the binary pulsar B0655+64, the stalwart control experiment for measurements of gravitational radiation damping in relativistic neutron-star binaries. Observed limits on evolution of the B0655+64 orbit provide new bounds on the existence of dipolar gravitational radiation, and hence on violation of the Strong Equivalence Principle. ",
        "introduction": "PSR B0655+64, in a highly circular one-day orbit with a $\\sim 0.8$ $M_{\\sun}$ white-dwarf companion, serves as a control experiment for measurements of orbital decay in the highly relativistic double-neutron-star binaries: General Relativity (GR) has predicted equally well the strong back-reaction to gravitational radiation for PSRs B1913+16 \\citep[][see also Weisberg \\& Taylor, this volume]{tay93b} and B1534+12 \\citep{1998ApJ...505..352S}, and the absence of detectable orbital evolution for PSR B0655+64 over two decades. The long-term stability of the B0655+64 orbit sets unique bounds on departures from GR that give rise to dipolar gravitational radiation \\citep{arz95,gol92}, the existence of which would represent a violation of the Strong Equivalence Principle (SEP), one of the basic tenets of GR. \\citet{2002PhRvD..66b4040G} examine the theoretical basis for dipolar gravitational radiation and suggest that bounds from binary pulsars may be competitive with future satellite experiments dedicated to probing SEP violation.  A definitive analysis of the available data for PSR B0655+64, and implications for a variety of alternative theories of gravitation, will be published separately (Arzoumanian et al.\\ 2003, in preparation). Following recent observations to extend our long-term monitoring program, I present here preliminary results on the orbital evolution of PSR B0655+64.  \\subsection{Binary Pulsars and the Strong Equivalence Principle}   The Strong Equivalence Principle posits that the response of a body to an external gravitational field is independent of the gravitational self-energy of the body; it would be violated if, for example, two objects with different gravitational binding energies were observed to fall at different rates. SEP is satisfied by postulate within GR; as a consequence, the lowest allowed multipole order of gravitational radiation is the ``electric'' quadrupole.  If SEP does not hold, however, emission of dipolar gravitational radiation is allowed, and lends itself in principle to indirect detection through decay of a binary orbit at a rate inconsistent with the quadrupole-order prediction of GR. Because self-gravitational effects are important in neutron stars, SEP can be tested with pulsar systems.  The most thoroughly studied theory of gravitation that violates SEP, the so-called ``Brans-Dicke'' scalar-tensor theory \\citep{bd61}, invokes the existence of a scalar gravitational field that couples to matter through a single dimensionless parameter, $\\omega_{B\\!D}$. The Brans-Dicke theory becomes indistinguishable from GR as $\\omega_{B\\!D}$ becomes arbitrarily large; solar-system experiments currently place a lower bound $\\omega_{B\\!D} \\ga 500$. Other alternatives to GR are also constrained by limits on dipolar gravitational radiation \\citep[see, e.g.,][]{de92}: recent observations of distant type Ia supernovae suggesting that the expansion of the universe is accelerating have prompted renewed interest in theories involving scalar fields (e.g., ``quintessence,'' \\citealp{1998Ap&SS.261..303C}) or long-range forces that may be responsible for the acceleration.  Such fields would radiate predominantly at dipole order.  The various contributions to observed changes in the orbital period $P_b$ of a binary system (neglecting tidal and mass-transfer effects) can be expressed as % \\begin{equation} \\label{brkdown.eq} \\left(\\frac{\\dot P_b}{P_b}\\right)_{\\rm obs} = \\left(\\frac{\\dot P_b}{P_b}\\right)_Q + \\left(\\frac{\\dot P_b}{P_b}\\right)_D + \\left(\\frac{\\dot P_b}{P_b}\\right)_{\\dot G} + \\left(\\frac{\\dot P_b}{P_b}\\right)_{\\bf a}, \\end{equation} where the subscripts $Q$, $D$, $\\dot G$, and $\\bf a$ denote the effects of quadrupolar and dipolar gravitational radiation, a change in the gravitational constant with time, and the varying Doppler shift from a relative acceleration of the solar system and pulsar binary\\footnote{To correct for Galactic accelerations and the ``Shklovskii'' effect, I follow \\citet{dt91}, using a new proper-motion measurement for B0655+64: $\\mu = 6.8\\pm1.1\\;{\\rm mas\\,yr}^{-1}$.}. In a large class of metric theories of gravitation, the radiation terms for a circular orbit are given by \\citep{ear75,wil81,gol92} \\begin{eqnarray} \\label{quad.eq} \\left(\\frac{\\dot P_b}{P_b}\\right)_Q &=& -\\frac{96}{5}{\\cal G}^{-4/3}\\left(\\frac{\\kappa_1}{12}\\right) n^{8/3}m_1m_2M^{-1/3}T_{\\sun}^{5/3} \\\\ \\label{dip.eq} \\left(\\frac{\\dot P_b}{P_b}\\right)_D &=& -\\kappa_D{\\cal G}^{-2}(s_1-s_2)^2n^2\\frac{m_1m_2}{M}T_{\\sun}, \\end{eqnarray} where $n \\equiv 2\\pi/P_b$, $m_1$ and $m_2$ are the pulsar and companion masses in solar units, $M \\equiv m_1 + m_2$, and $T_{\\sun}$ is the mass of the Sun in units of time. The dimensionless parameters $\\kappa_1$, $\\kappa_D$, and $\\cal G$ describe the strength of quadrupolar and dipolar radiation and the effective gravitational constant in a given theory. In GR, $\\kappa_1=12$, $\\kappa_D=0$, and ${\\cal G} = 1$; in the Brans-Dicke theory, $\\kappa_D{\\cal G}^{-2} = 2(2+\\omega_{B\\!D})^{-1}$. At the level of the current constraint on the dipole term $D$, the $\\dot G$ contribution to $\\dot P_b$ in Eq.~\\ref{brkdown.eq} can be neglected \\citep{arz95,ktr94}.  The quantities $s_1$ and $s_2$ in Eq.~\\ref{dip.eq} represent the ``sensitivities'' of the orbiting objects, the fractional change in binding energy of each star with $G$, \\begin{equation} \\label{sens.eq} s = -\\left(\\frac{\\partial \\ln m}{\\partial \\ln G}\\right)_N, \\end{equation} where $N$ is the total number of baryons. The sensitivity of white dwarfs is negligible; neutron stars are thought to have sensitivities in the range $0.15 < s < 0.40$ \\citep{wz89}, with larger values corresponding to ``softer'' equations of state. Results of X-ray burst oscillation modeling \\citep[e.g.,][]{2002ApJ...564..353N}, and the recent detection of atmospheric absorption lines from a neutron star \\citep{cottam02}, suggest that fairly stiff equations of state are appropriate. I therefore adopt a fiducial value $s = 0.2$ below. It is worth noting that the difference $(s_1 - s_2)^2$ in Eq.~\\ref{dip.eq} strongly suppresses dipolar radiation from NS-NS binaries like B1913+16, so that the celebrated agreement of the latter's orbital decay rate with the GR prediction does not usefully constrain the existence of gravitational radiation at dipole order. In scalar-tensor theories, gravitational binding energy takes on the role of an effective gravitational ``charge.'' Neutron star-white dwarf (NS-WD) binaries can therefore be expected to copiously emit dipolar radiation, if such radiation exists, because of the large difference in gravitational self-energy between the component stars.  \\subsection{PSR B0655+64} \\label{0655.sec}  PSR B0655+64 was discovered during a survey of the Northern sky made with the NRAO 300 Foot telescope \\citep{dbtb82}.  Timing observations began soon thereafter, and the companion was optically identified as a massive white dwarf by \\citet{kul86}. If we assume pulsar and companion masses of 1.4 $M_{\\sun}$ and 0.8 $M_{\\sun}$, Eq.~\\ref{quad.eq} predicts a rate of orbital decay within GR, \\( \\dot P_b^{\\rm GR} = -2\\times 10^{-14}, \\) corresponding to a decay timescale for the orbit of $\\tau_Q \\sim 150$~Gyr. ",
        "conclusions": "\\begin{table} \\begin{center} \\caption{\\label{nswd.tab} Relativistic NS-WD systems. Spin and orbital periods are shown alongside spin-down luminosity ($\\dot E$), orbital separation ($A$) and timescales for orbital evolution from quadrupolar ($\\tau_Q$) and dipolar ($\\tau_D$) radiation. The potential for heating and tidal interactions between the component stars increases rapidly for increasing $\\dot E$ and decreasing $A$.} \\begin{tabular}{lcccccc} \\tableline \\multicolumn{1}{c}{{PSR}} & {$P$} & {$P_b$} & {\\rule{0cm}{4mm}$\\log(\\dot E$[erg/s])} & {$A$} & {$\\tau_Q$} & {$\\tau_D$} \\\\ & {(ms)} & {(d)}\t& {} & {(lt-s)} & {(Gyr)} & {(Myr)} \\\\ \\tableline J1141$-$6545 & 393.9 & 0.20 & {33.4} & {\\phn4.5} & \\phn\\phn2 & \\phn2 \\\\ J0751+1807 & \\phn\\phn3.5 & 0.26 & {33.8} & {\\phn4.1} & \\phn20 & 15 \\\\ J1757$-$5322 & \\phn\\phn8.9 & 0.45 & {33.2} & {\\phn6.4} & \\phn20 & 15 \\\\ {B0655+64} & {195.6} & {1.03} & {30.6} & {11.3} & 150 & 60 \\\\ J1012+5307 & \\phn\\phn5.3 & 0.60 & {33.6} & {\\phn6.8} & 200 & 90 \\\\ J1435$-$6100 & \\phn\\phn9.3 & 1.35 & 33.1 & 14.2 & 250 & 90 \\\\ \\tableline \\tableline \\end{tabular} \\end{center} \\end{table}  Prospects for improving bounds on SEP violation and the existence of dipolar gravitational radiation are good. \\citet{dt92} show that, for constant data quality, measurement uncertainty for $\\dot P_b$ scales with data span $T$ as $T^{-5/2}$; moreover, data quality typically improves with time through improved instrumentation. Also, pulsar surveys continue to discover NS-WD systems similar to B0655+64 but with shorter orbital periods and millisecond pulse periods. Their higher rates of gravitational energy release coupled with higher timing precision suggest that these systems will surpass B0655+64 as laboratories for testing theories of gravitation% ---Table~\\ref{nswd.tab} lists the current sample of relativistic NS-WD systems. While B0655+64 remains a valuable object of further study, a potential caveat applies to relativistic studies of all close NS-WD binaries: very small orbital separations can give rise to non-relativistic interactions. Heating of the companion star by irradiation from the pulsar is thought to power tides and mass loss in a handful of low-mass binaries. Outflows and tides cause significant orbital torques (e.g., \\citealp{aft94}), which would overwhelm any small secular trend due to gravitational radiation. If indeed the closest NS-WD binaries (e.g., PSR B0751+18; see Nice et al., this volume) are ``clean'' systems, significant new constraints on SEP violation will emerge in the coming years."
    },
    "0212/astro-ph0212463_arXiv.txt": {
        "abstract": "A detailed quantitative analysis of the system $(ppe)$ placed in magnetic field ranging from $0 - 4.414 \\times 10^{13}\\,G$  is presented. The present study is focused on the question of the existence of the molecular ion $H_2^{+}$ in a magnetic field. As a tool, a variational method with an optimization of the form of the vector potential (optimal gauge fixing) is used.  It is shown that in the domain of applicability of the non-relativistic approximation the system $(ppe)$ in the Born-Oppenheimer approximation has a well-pronounced minimum in the total energy at a finite interproton distance for $B \\lesssim 10^{11}\\,G$, thus manifesting the existence of  $H_2^{+}$. For $B \\gtrsim 10^{11}\\,G$ and large inclinations (of the molecular axis with respect to the magnetic line) the minimum disappears and hence the molecular ion $H_2^{+}$ does not exist. It is shown that the most stable configuration of $H_2^{+}$ always corresponds to protons situated along the magnetic line. With magnetic field growth the ion $H_2^{+}$ becomes more and more tightly bound and compact, and the electronic distribution evolves from a two-peak to a one-peak pattern. The domain of inclinations where the $H_2^{+}$ ion exists reduces with magnetic field increase and finally becomes $0^o - 25^o$ at $B = 4.414 \\times 10^{13}\\,G$. Phase transition type behavior of variational parameters for some interproton distances related to the beginning of the chemical reaction $H_2^{+} \\leftrightarrow H + p$ is found. ",
        "introduction": "\\label{sec:intro}  Many years have passed since the moment when theoretical qualitative arguments were given that show that in the presence of a strong magnetic field the physics of atoms and molecules exhibits a wealth of new, unexpected phenomena even for the simplest systems \\cite{Kadomtsev:1971,Ruderman:1971}. In particular, a chance that unusual chemical compounds may be formed which do not exist without magnetic field was mentioned. In practice, the atmosphere of neutron stars, which is characterized by the presence of enormous magnetic fields $10^{12} - 10^{13}\\,G$, as well as other astronomical objects carrying large magnetic fields ($>10^{8}\\,G$) provide a valuable paradigm where this physics could be realized. Recently, the experimental data collected by the {\\it Chandra} X-ray observatory revealed certain irregularities in the spectrum of an isolated neutron star 1E1207.4-5209. These irregularities can be interpreted as absorption features at $\\sim 0.7\\,$ KeV and $\\sim 1.4\\,$ KeV of possible atomic or molecular nature \\cite{Sandal:2002}.  One of the first general features observed in standard atomic and molecular systems placed in a strong magnetic field is an increase of both total and binding energies, accompanied by a drastic shrinking of the electron localization length in both the longitudinal and transverse directions. Naturally, this leads to a decrease of the equilibrium distance with magnetic field growth. This behavior can be considered to be a consequence of the fact that for large magnetic fields the electron cloud takes a needle-like form extended along the magnetic field direction and the system becomes effectively quasi-one-dimensional \\cite{Ruderman:1971}. It is obvious that the phenomenon of quasi-one-dimensionality enhances the stability of standard atomic and molecular systems from the electrostatic point of view. In particular, molecules become elongated along the magnetic line forming a type of linear molecular polymer (for details see the review articles \\cite{Garstang:1977, Lai:2001}). It also hints at the occurrence of exotic atomic and molecular systems which do not exist in the absence of a magnetic field. Motivated by these simple observations it was shown in Refs.\\cite{Turbiner:1999, Lopez-Tur:2000} that three and even four protons can be bound by one electron. This shows that exotic one-electron molecular systems $H_3^{2+}$ and $H_4^{3+}$ can exist in sufficiently strong magnetic fields in the form of linear polymers. However, the situation becomes much less clear (and also much less studied) when the nuclei are not aligned with the magnetic field direction, and thus in general do not form  a linear system. Obviously, such a study would be important for understanding the kinetics of a gas of molecules in the presence of a strong magnetic field. As a first step towards such a study, even the simplest molecules in different spatial configurations deserve attention. Recently, a certain spatial configuration of $H_3^{2+}$ was studied in detail \\cite{Lopez-Tur:2002}. It was shown that in the range of magnetic fields $10^8 < B < 10^{11}\\,G$ the system $(pppe)$, with the protons forming an equilateral triangle perpendicular to the magnetic lines, has a well-pronounced minimum in the total energy for a certain size of triangle. The goal of the present work is to attempt for the first time to carry out an extensive quantitative investigation of the ground state of $H_2^+$ in the framework of a single approach in its entire complexity: a wide range of magnetic field strengths ($0 - 4.414 \\times 10^{13}\\,G$), arbitrary (but fixed) orientation of the molecular axis with respect to the magnetic line and arbitrary internuclear distances. We are going to carry out this study in the Born-Oppenheimer approximation at zero order -- assuming protons to be infinitely heavy charged centers. In principle, when the molecular axis is perpendicular to the magnetic line the system $(ppe)$ acquires extra stability from the electrostatic point of view. Electrostatic repulsion of the classical protons is compensated for by the Lorentz force acting on them.  It is well known that the molecular ion $H_2^+$ is the most stable one-electron molecular system in the absence of a magnetic field. It remains so in the presence of a constant magnetic field unless $B \\gtrsim 10^{13}\\,G$, where the exotic ion $H_3^{2+}$ appears to be the most bound (see \\cite{Lopez-Tur:2000}). The ion $H_2^+$ has been widely studied, both with and without the presence of a magnetic field, due to its importance in astrophysics, atomic and molecular physics, solid state and plasma physics (see \\cite{Garstang:1977}-\\cite{Potekhin} and references therein). The majority of the previous studies were focused on the case of the parallel configuration, where the angle between the molecular axis and the magnetic field direction is zero, $\\tha=0^o$. The only exception is Ref.\\cite{Schmelcher}, where a detailed quantitative analysis was performed for any $\\tha$ but for a single magnetic field $B=1\\,a.u.\\,$. Previous studies were based on various numerical techniques, but the overwhelming majority used different versions of the variational method, including the Thomas-Fermi approach. As a rule, in these studies the nuclear motion was separated from the electronic motion using the Born-Oppenheimer approximation at zero order (see above). It was observed at the quantitative level that magnetic field growth is always accompanied by an increase in the total and binding energies, as well as a shrinking of the equilibrium distance. As a consequence it led to a striking conclusion about the drastic increase in the probability of nuclear fusion for $H_2^+$ in the presence of a strong magnetic field \\cite{Kher}.  In the present study we will also use a variational method. Our consideration will be limited by a study of the $1_g$-state, which realizes the ground state of the system if the bound state exists \\footnote{After a straightforward separation of the spin part of wavefunction, the original Schroedinger equation becomes a scalar Schroedinger equation. Then it can be stated that a nodeless eigenfunction corresponds to the ground state (Perron theorem)}. We will construct state-of-the-art, non-straightforward, `adequate' trial functions consistent with a variationally optimized choice of vector potential. We should stress that a proper choice of the form of the vector potential is one of the crucial points which guarantee the adequacy and reliability of the consideration. In particular, a proper position of the gauge origin, where the vector potential vanishes, is drastically important, especially for large interproton distances. For the parallel configuration, $\\tha=0^o$ the present work can be considered as an extension (and also an development) of our previous work \\cite{Lopez:1997}. It is necessary to emphasize that we encounter several new physical phenomena which occur when the molecular axis deviates from the magnetic field direction. If the magnetic field is sufficiently strong, $B\\gtrsim 10^{11}\\,G$ and inclination $\\tha$ is larger than a certain critical angle the ion $H_2^+$ does not exist in the contrary to a prediction in Refs.\\cite{Larsen}, \\cite{Kher}, \\cite{Turbiner:2001}. This prediction was based on an improper gauge dependence of the trial functions, which caused a significant loss of accuracy and finally led to a qualitatively incorrect result. We find that in the weak field regime the $(ppe)$ system in the equilibrium position at any inclination, the electronic distribution peaks at the positions of the protons. While at large magnetic fields the electronic distribution is characterized by single peak at the midpoint between two protons. This change from a two-peak to a one-peak configuration appears around $B \\sim 10^{10}-10^{11}\\,G$ with a slight dependence on the inclination angle $\\tha$. From a physical point of view the former means that the electron prefers to stay in the vicinity of a proton. This can be interpreted as a dominance of the $H$-atom plus proton configuration. The latter situation implies that the electron is `shared' by both protons and hence such a separation to $H$-atom plus proton is irrelevant. Therefore, we can call the two-peak situation `ionic' coupling, while the one-peak case can be designated as `covalent' coupling, although this definition differs from that widely accepted in textbooks (see, for example \\cite{LL}). Thus, we can conclude that a new phenomenon appears - as the magnetic field grows the type of coupling changes from `ionic' to `covalent'. At large internuclear distances the electron is always attached to one of the charged centers, so the coupling is `ionic'.  One particular goal of our study is to investigate a process of dissociation of the $(ppe)$ system: $H_2^+ \\rar H + p$ which appears with increase of interproton distance. It is clear from a physical point of view that at large distances the electronic distribution should be first of the two-peak type and then should change at asymptotically large distances, to a single-peak one, but with a peak at the position of one of the protons. Somehow this process breaks permutation symmetry and we are not aware of any attempt to describe it. In our analysis this phenomenon appears as a consequence of a change of position of the gauge origin with increase of interproton distance.  From the physical point of view it is quite interesting to note how the $(ppe)$ system behaves at very large interproton distances. This domain is modelled by an $H$-atom plus proton interaction. The interaction corresponds to (magnetic-field-inspired-quadrupole) + charge interaction and is dominant comparing to the standard Van der Waals force. For small inclinations the above interaction is attractive as in the Van der Waals case, but becomes {\\it repulsive} for large inclinations. This implies that the potential curves approach to asymptotic value of total energy at the large interproton distances from above contrary to the Van der Waals case.  The Hamiltonian which describes two infinitely heavy protons and one electron placed in a uniform constant magnetic field directed along the $z-$axis, ${\\bf B}=(0,0,B)$ is given by (see e.g. \\cite{LL}) \\begin{equation} \\label{Ham} {\\cal H} = {\\hat p}^2 + \\frac{2}{R} -\\frac{2}{r_1} -\\frac{2}{r_2} + ({\\hat p} {\\cal A}+{\\cal A}{\\hat p}) +  {\\cal A}^2 \\ , \\end{equation} (see Fig.1 for notations), where ${\\hat p}=-i \\na$ is the momentum, ${\\cal A}$ is a vector potential, which corresponds to the magnetic field $\\bf B$. Hence, the total energy $E_T $ of $H_2^+$ is defined as the total electronic energy plus the Coulomb energy of proton repulsion. The binding energy is defined as an affinity to have the electron at infinity, $E_b=B-E_T$. The dissociation energy is defined as affinity to have a proton at infinity, $E_d=E_H-E_T$, where $E_H$ is the total energy of the hydrogen atom in magnetic field $B$.   \\begin{figure}[tb] \\begin{center} \\includegraphics*[width=4.5in]{h2+incl.ps} \\caption{Geometrical setting  for the $H_2^{+}$ ion placed in a magnetic field directed along the $z$-axis. The protons are situated in the $y-z$ plane at a distance $R$ from each other and marked by bullets. $O$ is the origin of coordinates which is chosen to be on the bold-dashed line which connects protons; $O'(0,Y,Z)$ is the mid-point between protons. It is assumed that the gauge center coincides with $O$. $OO'$ measures the distance between the gauge center and the mid-point between the proton positions (see text and eq.(4)).} \\label{fig:1} \\end{center} \\end{figure}  Atomic units are used throughout ($\\hbar$=$m_e$=$e$=1) albeit energies are expressed in Rydbergs (Ry). Sometimes, the magnetic field $B$ is given in a.u. with $B_0= 2.35 \\times 10^9\\,G$ \\footnote{In absence of convention, some results presented in literature are obtained for $B_0= 2.3505 \\times 10^9\\,G$. Thus, making a comparison of the results obtained by different authors, this fact should be taken into account}. ",
        "conclusions": "We have carried out an accurate, non-relativistic calculation in the Born-Oppenheimer approximation for the lowest state of the $H_2^+$ molecular ion for different orientations of the magnetic field with respect to the molecular axis. We studied constant uniform magnetic fields ranging from zero up to $B = 4.414 \\times 10^{13}\\,G$, where non-relativistic consideration holds.  For all magnetic fields studied there exist a region of inclinations for which a well-pronounced minimum in the total energy surface for the $1_g$ state of the system $(ppe)$ is found. This shows the existence of the $H_2^+$ molecular ion for magnetic fields $B = 0 - 4.414 \\times 10^{13}\\,G$. The smallest total energy is always found to correspond to the parallel configuration, $\\tha=0^{\\circ}$, where protons are situated along the magnetic line. The total energy increases, while the binding energy decreases monotonically as the inclination angle grows. The rate of total energy increase as well as binding energy decrease is seen to be always maximal for the parallel configuration. The equilibrium distance exhibits quite natural behavior as a function of the orientation angle $\\tha$ -- for fixed magnetic field the shorter equilibrium distance always corresponds to the larger $\\tha$.  Confirming the qualitative observations made by Khersonskij \\cite{Kher} for the $1_g$ state in the contrast to statements in \\cite{Larsen, Wille:1988}, we demonstrate accurately that the $H_2^+$-ion does not exist at a certain range of orientations for magnetic fields $B \\gtrsim 2 \\times 10^{11}\\,G$. As the magnetic field increases the region of inclinations where $H_2^+$ does not exist is seen to broaden, reaching rather large domain $25^{\\circ}\\lesssim \\tha \\leqslant 90^{\\circ}$ for $B=4.414 \\times 10^{13}\\,G$.  We find that the electronic distributions for $H_2^+$ in the equilibrium position are qualitatively different for weak and large magnetic fields. In the domain $B < 10^{10}\\,G$ the electronic distribution for any inclination has a two-peak form, peaking near the position of each proton. On the contrary for $B>10^{11}\\,G$ the electronic distribution always has a single peak form with the peak near the midpoint between the protons for any inclination. This implies a physically different structure for the ground state - for weak fields the ground state can be modelled as a `superposition' of hydrogen atom and proton, while for strong fields such modelling is not appropriate.  Unlike standard potential curves for molecular systems in the field-free case, we observe for $\\tha > 0^o$ that each curve has a maximum and approaches to the asymptotics at $R \\rar \\infty$ from above. The electronic distribution evolves with $R$ from a one-peak form at small $R$ to a two-peak one at large $R$. There exists a certain critical $R_{cr}$ at which one of peaks starts to diminish, manifesting a breaking of permutation symmetry between the protons and simultaneously the beginning of the chemical reaction $H_2^+ \\rar H + p$.  Combining all the above-mentioned observations we conclude that for magnetic fields of the order of magnitude $B \\sim 10^{11}\\,G$ some qualitative changes in the behavior of the $H_2^+$ ion take place. The behavior of the variational parameters also favors this conclusion. This hints at the appearance of a new scale in the problem. It might be interpreted as a signal of a transition to the domain of developed quantum chaos (see, for example, \\cite{chaos})."
    },
    "0212/astro-ph0212324_arXiv.txt": {
        "abstract": "Research on interstellar \\dammo\\ is reviewed and updated from the discovery papers. Results from observations of a dozen sources at centimeter and submillimeter wavelengths are presented. The two data sets are consistent, but do not constrain the excitation conditions in the \\dammo-emitting gas. The column density ratios of \\ammo/\\nhhd, \\nhhd/\\nddh\\ and \\dammo/\\nddh, observed in similar sized beams, are $\\approx$10, and present a challenge for both the gas-phase and the grain-surface chemistry scenarios of deuterium fractionation.  The role of shocks and of CO depletion in deuterium chemistry is discussed. ",
        "introduction": "\\label{sec:intro}  Deuterium-bearing molecules are important as probes of the very cold phases of molecular clouds, prior to star formation. Moreover, the isotopic composition of molecules is a valuable clue to their formation mechanism.  In the gas phase, deuterium fractionation occurs by reactions with \\hhdp, CH$_2$D$^+$ and C$_2$HD$^+$ (Millar, this volume).  Alternatively, reactions on the surfaces of dust grains may enhance molecular D/H ratios. Both methods require low temperatures ($<$20~K) and high densities ($>10^5$~\\ccm) such as found in dark clouds and pre-stellar cores. The grain surface route also needs a mechanism to return the molecules to the gas phase where they can be observed by the means of submillimeter spectroscopy. The gas-phase route can account for the deuteration of HCO$^+$ and N$_2$H$^+$, while grain surface chemistry may be important for \\hhco\\ and \\chhhoh.  Although astronomical detections of both singly and doubly deuterated \\ammo\\ exist (\\cite{roueff00}), the deuteration mechanism for \\ammo\\ is not clear.  Despite early claims, searches for \\ammo\\ ice based on the 2.21~and 9.0~\\mic\\ features have not been successful (\\cite{taban02}), but the 3.47~\\mic\\ feature indicates abundances of up to 7\\% of \\hho\\ ice (\\cite{dart02}). However, existence in the solid state does not imply formation on grain surfaces, since \\ammo\\ may just passively accrete after forming in the gas phase.  Soon after \\cite*{rodg01} proposed \\dammo\\ as test of theories of interstellar deuterium chemistry, searches were initiated.  The rotational energy levels of symmetric top molecules like \\dammo\\ are labeled by the total angular momentum $J$ and its projection on the molecular symmetry axis $K$. Inversion splits these levels into states which are symmetric (\\textit{s}) and antisymmetric (\\textit{a}) upon reflection in the plane of the D~atoms. Coupling of the $^{14}$N nuclear spin $I$ with the rotational angular momentum $J$ causes additional hyperfine splitting ($F=I+J$).  Current instrumentation cannot resolve the even finer splitting due to D, which, owing to its small quadrupole moment, is only $\\approx$200~kHz.  A complete line list is available from the CDMS catalog (\\cite{mull01}) at the URL {\\tt http://www.cdms.de}.  Under the conditions favourable for deuterium enhancement, the strongest lines of \\dammo\\ should be the ground state rotational line ($1_0^a\\to0_0^s$) at 309.9~GHz and the (1,1) inversion line ($1_1^a\\to1_1^s$) at 1589~MHz, which are both accessible from the ground.  Because the Pauli exclusion principle only holds for fermions, certain states are forbidden for \\ammo\\ but allowed in \\dammo. For example, the \\dammo\\ ground state rotational line also has a $1_0^s\\to0_0^a$ component at 306.7~GHz, which does not exist for \\ammo. However, spin statistics make this line 10 times weaker than the $1_0^a\\to0_0^s$ line, which is why it has escaped detection so far. ",
        "conclusions": ""
    },
    "0212/astro-ph0212438_arXiv.txt": {
        "abstract": "s{ Core-collapse supernovae can be used to place limits on dark matter candidate particles, but the strength of these limits depends on the depth of our theoretical understanding of these astrophysical events. To date, limitations on computing power have prevented inclusion of all the physics that would constitute a realistic simulation. The TeraScale Supernova Initiative (TSI) will overcome these obstacles in the next few years, elucidating the explosion mechanism and other phenomena closely associated with the core collapse of massive stars. } ",
        "introduction": " ",
        "conclusions": "While many simulations have been performed over the years, it cannot yet be claimed that the supernova explosion mechanism is understood. Models with energy- and angle-dependent neutrino transport have been studied in spherical symmetry, and explosions have not been seen. Many (though not all) models with multidimensional hydrodynamics do exhibit explosions, but these have employed neutrino transport that is too crude to make firm conclusions about the explosion mechanism. Computing power has advanced to the point that models with {\\em both} sophisticated neutrino transport and multidimensional hydrodynamics are within reach; the TeraScale Supernova Initiative (TSI) is one effort underway to perform such simulations.  Finally, in conclusion, a word on the subject of the conference: dark matter. The nascent neutron star is a hot (temperature of order 50 MeV) and dense (baryon mass density of order $10^{14}$ g cm$^{-3}$) environment. Should hypothetical particles like the axion (a cold dark matter candidate) or a keV-mass sterile neutrino (a warm dark matter candidate) exist, they could be produced in the extreme conditions present in the newly born neutron star. However, copious production and emission of such weakly coupled particles would cause the neutron star to cool more quickly than it would if neutrinos were fastest means of cooling. The handful of neutrinos detected from supernova SN1987A confirms the basic theoretical understanding of stellar core collapse, with associated trapping of neutrinos and their subsequent emission on a time scale of several seconds; this allows limits to be placed on the coupling strength of hypothetical dark matter particles like the axion\\cite{raffelt99} and sterile neutrino.\\cite{abazajian01} Presently, these limits are of necessity rather conservative, due to a lack of detailed understanding of the explosion mechanism and subsequent uncertainties about precise conditions in the nascent neutron star. A new generation of simulations---such as those being pursued by TSI---promises to reveal the explosion mechanism and paint a detailed and realistic picture of the physical conditions in the hot and dense neutron star, providing a basis for strengthened limits on the properties of dark matter candidate particles."
    },
    "0212/astro-ph0212112_arXiv.txt": {
        "abstract": "{ The propagation of Galactic Cosmic Ray nuclei having energies between 100 MeV/nuc and several PeV/nuc is strongly believed to be of diffusive nature. The particles emitted by a source located in the disk do not pervade the whole Galaxy, but are rather confined to a smaller region whose spatial extension is related to the height of the diffusive halo, the Galactic wind and the spallation rate. Following the pioneering work of \\cite{Jones78}, this paper presents a general study on the {\\em spatial} origin of cosmic rays, with a particular attention to the role of spallations and Galactic wind. This question is different, and to a certain extent disconnected, from that of the {\\em origin} of cosmic rays. We find the regions of the disk from which a given fraction of cosmic rays detected in the Solar neighborhood were emitted ($f$-surfaces). After a general study, we apply the results to a realistic source distribution, with the propagation parameters obtained in our previous systematic analysis of the observed secondary-to-primary ratios \\citep{Maurin02b}. The shape and size of these $f$-surfaces depend on the species as well as on the values of the propagation parameters. For some of the models preferred by our previous analysis (i.e. large diffusion slope $\\delta$), these $f$-surfaces are small and in some extreme cases only a fraction of a percent of the whole Galactic sources actually contribute to the Solar neighborhood Cosmic Ray flux. Moreover, a very small number of sources may be responsible for more than 15 \\% of the flux detected in the Solar neighborhood. This may point towards the necessity to go beyond the approximations of both homogeneity and stationarity. Finally, the observed primary composition is dominated by sources within a few kpc. ",
        "introduction": "The propagation of charged Cosmic Ray nuclei, in the energy range going from a few 100 MeV/nuc and a few PeV/nuc, is strongly affected by the Galactic magnetic field. It is a diffusive process, so that the Cosmic Rays emitted by a single source  spread out in time, pervade the whole Galaxy, and can escape the Galaxy when reaching its boundaries. Those coming from a source located far from the Sun have a larger probability of escaping than reaching the Solar neighborhood. It is the opposite for nearby sources, so that the cosmic ray fluxes in the solar neighborhood are more sensitive to the properties of the local sources (as opposed to the remote sources). Other effects like spallations and Galactic wind further limit the distance Cosmic Rays travel before being detected. Some consequences of the Galactic wind were studied in \\citet{Jones78} where convective escape was compared to escape through the top and bottom boundaries of the Galaxy.  The goal of this paper is to go one step beyond by providing a general study on the {\\em spatial} origin of cosmic rays, i.e. to answer the question \"from which region of the Galaxy were emitted the cosmic rays detected in the solar neighborhood?\". This question is different, and to a certain extent disconnected, from that of the {\\em origin} of cosmic rays (\"What are the astrophysical objects which are responsible for the acceleration of cosmic rays?\") which is still much debated. We believe that it is nevertheless an interesting question, for several reasons. First, we find that the answer may cast some doubt on the validity of the stationary model, upon which most studies on Cosmic Rays are based. Second, it gives some clues about the spatial range beyond which the cosmic ray studies are blind to the sources. Finally, this study may be of interest to optimize the propagation codes based on Monte-Carlo methods, by focusing the numerical effort on the sources that really contribute to the detected flux.  The reader who does not want to go through the pedagogical progression can go directly from the general presentation of the method in $\\S$~\\ref{general_formula} to its application in realistic cases in $\\S$~\\ref{BC_induced}. For the others, the effect of escape is studied in $\\S$~\\ref{bound} and that of spallations and Galactic wind is studied in $\\S$~\\ref{spal_et_conv}. Then, $\\S$~\\ref{resultats_de_base} studies the effect of a realistic source distribution. Finally, the fully realistic case is considered in $\\S$~\\ref{BC_induced}. The results and the perspectives are discussed in the last section. For convenience, we will use the word $f$-surfaces to describe the surfaces in the thin disk within which the sources form the fraction $f$ of cosmic rays detected at the observer location. ",
        "conclusions": "The question of the source distribution is very present in Cosmic Ray physics. With the occurrence of the old problem of short pathlengths distribution in leaky box models (see for example \\citet{webber98}), \\cite{lezniak79} studied a diffusion model with no-near source in the Solar neighborhood. Later, \\cite{webber93a, webber93b} propagated $\\delta$-like sources with diffusion generated by a Monte Carlo random walk for the same purpose. \\cite{Brunetti00} follow this line but they introduce random walks in a more realistic environment, i.e. circular, elliptical and spiral magnetic field configurations. In a more formal context, \\cite{lee79} used a statistical treatment of means and fluctuations (see references therein) to characterize amongst others the possibility that nearby recent sources may dominate the flux of primaries. Finally, it is known that the present Cosmic Ray models are not able to reproduce accurately for example proton-induced $\\gamma$-rays measurements. To illustrate this point, we plot in Fig.~\\ref{fig:dist_gamma} the radial distribution of protons obtained with the same diffusion parameters as used above. None of the models shown match the data. \\begin{figure}[ht] \\centerline{ \\includegraphics[width=0.85\\columnwidth]{ms3376_f16.eps} } \\caption{Radial distribution of the proton flux for the models discussed in this study, compared to the source radial distribution of \\citet{case98} given Eq.~(\\ref{case_batth}). For each of the values $L=2$~kpc and $L=10$~kpc, the three values $\\delta=0.35$, 0.6 and 0.85 are presented, the flatter distribution corresponding to the lower $\\delta$. Also shown is the gamma-ray emissivity per gas atom ({\\sc cos-B} \\citet{bloemen89}, which is proportional to the proton flux, as given by {\\sc cos-B} (open circles, \\citet{bloemen89}) and {\\sc egret} (triangles, \\citet{strong_mattox96}), along with the proton flux obtained with the \\citet{Strong98} distribution (see Sec.~\\ref{resultats_de_base}).} \\label{fig:dist_gamma} \\end{figure} One is left with two alternatives: either modifying the source distribution (for example, the distribution of \\citet{Strong98} yields a better agreement), or giving up the assumption that the diffusion parameters apply to the Galaxy as a whole \\citep{breit02}. It is thus of importance to understand to what extent the Cosmic Rays detected on Earth are representative of the distribution of the sources in the whole Galaxy.  We provide an answer to this question under the two important hypotheses that the source distribution is continuous and that we have reached a stationary regime: most of the Cosmic Rays that reach the Solar neighborhood were emitted from sources located in a rather small region of the Galactic disk, centered on our position. The quantitative meaning of ``rather small\" depends on the species as well as on the values of the diffusion parameters. For the generic values $\\delta = 0.6$ and $L=6$ kpc chosen among the preferred values fitting B/C (see Sec.~\\ref{BC_induced}), half of the protons come from sources nearer than 2 kpc, while half of the Fe nuclei come from sources nearer than 500 pc. Another way to present this result is to say that the fraction of the whole Galactic source distribution that actually contributes to the Solar neighborhood Cosmic Ray flux can be rather small. For the generic model just considered, 8 \\% (resp. 1.5 \\%) of the sources are required to account for 90 \\% of the proton (resp. Fe) flux. These fractions are smaller for higher $\\delta$ and smaller $L$. To summarize, the observed Cosmic Ray primary composition may be dominated by sources within a few kpc, so that a particular care should be taken to model these source, spatially as well as temporally \\citep{Maurin03b}.  Independently of all the results, this study could be used as a check for more sophisticated Monte Carlo simulations that will certainly be developed in the future to explore inhomogeneous situations. Several other consequences deserve attention. First, the results may point towards the necessity to go beyond the approximations of both continuity and stationarity. In particular, it could be that only a dynamical model, with an accurate spatio-temporal description of the nearby sources, provides a correct framework to understand the  propagation of Galactic Cosmic Rays. The contribution from nearby sources would be very different in the low energy ($\\sim$~GeV/nuc) or in the high energy regime ($\\sim$~PeV) compared to the stationary background. Second, as discussed in Sec.~\\ref{progenitors}, the diffusion parameters derived from the observed B/C ratio have only a local validity, and one should be careful before applying them to the whole Galaxy, since the Cosmic Rays are blind to most of it."
    },
    "0212/astro-ph0212262_arXiv.txt": {
        "abstract": "{We have measured the Hubble constant $H_0$ in NGC~564 at $cz \\approx 5800$ km~s$^{-1}$ and in NGC~7619 at $cz \\approx 3700$ km s$^{-1}$ with deep I--band Surface Brightness Fluctuation distance measurements at the ESO Very Large Telescope (VLT). We obtain~$H_0 = 70 \\pm 7 \\pm 6$ km~s$^{-1}$/Mpc for NGC~564 and $H_0 = 68 \\pm 6 \\pm 6$ for NGC~7619. The actual SBF sample used for the measurement of $H_0$ in the Hubble Space Telescope Key Project on the Extragalactic Distance Scale (Freedman et al. 2001) amounts to six galaxies.  When we combine the measurements from this work with our previous VLT I--band SBF distance measurement in IC~4296 (Mei et al. 2000), we obtain~:~$H_0$ = 68 $\\pm$ 5 $\\pm$ 6 km~s$^{-1}$/Mpc. When we add the Freedman et al. (2001) SBF sample, we obtain $H_0 = 71 \\pm 4 \\pm 6$ km s$^{-1}$/Mpc. ",
        "introduction": "Future Cosmic Microwave Background missions like MAP (http://map.gsfc.nasa.gov/) and Planck (http://astro.estec.esa.nl/SA-general/Projects/Planck/) promise to obtain definitive measurements for cosmological parameters, achieving an accuracy around $1\\%$. Over the last few years, however, the accuracy of such measurements has improved considerably, as many efforts have been carried out to obtain accurate measurements for the Hubble constant $H_0$, to reduce the uncertainty in its determination from 50$\\%$ to approximately $10\\%$. The most important among these efforts has been the Hubble Space Telescope Key Project on the Extragalactic Distance Scale (HSTKP; Freedman et al. 2001 and references therein) where the value of $H_0$ was measured to be $H_0 = 72 \\pm 8$  km~s$^{-1}$, based on a coherent Cepheid calibration of four of the most dependable and precise methods for galaxy distance measurements: Tully--Fisher, Fundamental Plane, SNIa and Surface Brightness Fluctuations. Other important contributions were provided, among others, by Ajhar et al. (2001); Blakeslee et al. (2001); Jensen et al. (2001); Gibson \\& Stetson (2001); Liu \\& Graham (2001); Tonry et al. (2000);  Suntzeff et al. (1999); Jha et al. (1999); Giovanelli et al. (1997); Hjorth \\& Tanvir (1997).  $H_0$ measurements are affected not only by uncertainties in the derivation of the cosmological distance ladder, but also by the presence of local deviations from the smooth Hubble flow. The average peculiar velocity for galaxies and clusters of galaxies are of the order of 100 - 300~km~s$^{-1}$ (Davis et al. 1997; Giovanelli et al. 1998; Dale et al. 1999). Therefore to minimise their effect on the determination of $H_0$, it is preferable to select target objects beyond 4000~ km~s$^{-1}$, where the peculiar motions amount to no more than $\\approx 10 \\% - 15 \\% $ of the total velocity. This requirement has led to a number of projects trying to build samples of redshift independent distance estimates restricted to galaxies located beyond 4000 - 5000~km~s$^{-1}$ (Dale et al. 1999; Hudson et al. 2001; Colless et al. 2001; Jensen et al. 2001).  Among the methods used to obtain accurate galaxy distance estimates, Surface Brightness Fluctuations (SBF) are, at present, one of the most accurate and better physically understood methods available.  The method was introduced by Tonry \\& Schneider 1988 (a recent review has been given by Blakeslee et al. 1999a) and is based on the very simple fact that the Poissonian distribution of unresolved stars in a galaxy produces fluctuations in each pixel of the galaxy image.  The variance of these fluctuations is inversely proportional to the square of the galaxy distance. The SBF amplitude is defined to be precisely this variance, normalised to the mean flux of the galaxy in each pixel \\cite{ts88}.   \\begin{table*}[!htf]  \\begin{flushleft} \\begin{tabular} {|c|c|c|c|c|c|c|c|c|} \\hline Name&RA &DEC &Type&$m_V^1$&$V-I^1$&$cz_{LG}^2 $&$cz_{CMB}^2$&$V^{pec 2}_{CMB}$ \\\\ \\hline &(2000)&(2000)&&mag&mag&km s$^{-1}$& km s$^{-1}$&km s$^{-1}$ \\\\ \\hline NGC 564&01 27 48.31 & -01 52 47.68 &E&13.01&1.2&5843& 5037&-52 \\\\ \\hline NGC 7619&23 20 14.68 &+08 12 23.3 &E&11.90&1.23&3758&3519&-256 \\\\ \\hline \\end{tabular}  \\end{flushleft}  $^1$ Prugniel $\\&$ Heraudeau 1998, Poulain $\\&$ Nieto 1994 \\\\ $^2$ Kelson et al. 2000 and references therein, Giovanelli 1998, Dale 1999  \\caption{General properties of NGC 564 and NGC 7619, from the web archive Simbad (http://simbad.u-strasbg.fr) and the NASA/IPAC Extragalactic Database (NED) (http://nedwww.ipac.caltech.edu/)} \\label{tab-vlt} \\end{table*}   \\begin{table*}[!htf]  \\begin{flushleft} \\begin{tabular} {|c|c|c|} \\hline Port&FORS1 Ron ($e^{-}$)&Gain ($e^{-}$/ADU)\\\\ \\hline  A&5.01$\\pm$0.16&1.48$\\pm$0.04 \\\\ B&5.42$\\pm$0.17&1.78$\\pm$0.05 \\\\ C&5.18$\\pm$0.16&1.56$\\pm$0.04 \\\\ D&5.03$\\pm$0.16&1.63$\\pm$0.05 \\\\ \\hline  \\end{tabular}  \\end{flushleft}   \\caption{Read--out Noise and Gain for the four ports of FORS1 on VLT UT1} \\label{tab-fors} \\end{table*}    The absolute magnitude of the fluctuation is not a constant, but depends on the age and metallicity of the stellar population within the galaxy. Tonry et al. (1997) and Tonry et al. (2001) have quantified these dependencies using an extensive sample of I and V band observations, which are used to empirically calibrate the dependence of the I -- band fluctuation magnitude on (V -- I) colour. The absolute calibration is then obtained using a set of 5 galaxies with independent Cepheid distances.  As part of the HSTKP, Ferrarese et al. (2000a) have obtained their own calibration based on accurate Cepheid distances to six spiral galaxies with SBF measurements within 1200~km~s$^{-1}$. The Cepheid distances were part of a larger dataset used by HSTKP. The two calibrations are consistent within the errors.  Alternatively, theoretical calibrations have been obtained starting from synthesis models of the galaxy stellar population (Ajhar et al. 2001; Liu et al. 2000; Blakeslee et al. 2001). From Ajhar et al. (2001) the zero points of these calibrations provide an average value of the Hubble constant ($H_0 = 77 \\pm 7$~km~s$^{-1}$/Mpc) that is in good agreement with the observational results.  With HST and 8-m class telescopes SBF measurements have been extended beyond the local universe ($cz < 4000$~km~s$^{-1}$). While the survey by Tonry et al. (2001) was limited to galaxies with $cz < 4000$~km s$^{-1}$, using HST observations it has been possible to determine distances for more distant galaxies (Thomsen et al. 1997; Lauer et al. 1998; Pahre et al. 1999; Jensen et al. 2001; Liu et al. 2001; Ajhar et al. 2001), up to approximately $10,000$~km~s$^{-1}$.  In this paper we extend our ground-based I-band SBF measurements to beyond $\\approx 5000$~km~s$^{-1}$. Advantages and disadvantages of I- versus K-band ground based SBF observations were discussed using theoretical simulations in Mei et al. (2001).  Deep I--band SBF observations were obtained for two galaxies : NGC~564, in Abell 194, at $cz \\approx 5800$~km~s$^{-1}$, and NGC~7619 in the Pegasus cluster at $cz \\approx 3700$~km~s$^{-1}$.  Both Abell 194 and Pegasus are fiducial clusters in the derivation of observational templates for the Tully--Fisher (TF) and Fundamental Plane (FP) relations (see for example Mould et al. 1991, 1993; Jorgensen et al. 1996; Giovanelli et al. 1998; Scodeggio et al. 1998; Dale et al. 1999; Hudson et al. 2001), while Abell 194 is also part of the sample of clusters studied by Lauer \\& Postman (1992, 1994). Although distance estimates obtained with the TF and FP relations are comparatively less accurate than those obtained with the SBF method, they are relatively easier to obtain, even for objects located at larger distances than those sampled so far by SBF observations. Therefore the most complete and farthest reaching mappings of the local universe currently make use of TF- and FP-based distances, and they will likely do so in the near future. One of the main motivations behind the selection of targets for this work was the possibility of providing a bridge between SBF distance estimates and TF or FP results, as this is the most direct way to improve the accuracy of the calibration for the latter relations, and consequently that of all distance estimates based on these relations. This effort extends to larger distances the matching of SBF and FP distance scales recently presented by Blakeslee et al. (2002).  In $\\S$ 2 we describe our observations, and the SBF analysis.  The $H_0$ measurements and a summary of this analysis are presented in $\\S$ 3 and 4.     \\begin{table*}[!htf]  \\begin{flushleft} \\begin{tabular} {ccccccc} \\hline  Galaxy&Band&Date&Frames&Exp.Time (sec)\\\\ \\hline NGC 7619&V& 3 Nov 00&5&100\\\\ &I&5 Oct 00&4&60 \\\\ &I&5 Oct 00&55&150 \\\\ NGC 564&V&19 Aug 01&5& 100\\\\  &I&3 Nov 00&81&65\\\\ &I&2 Dec 00&53&65\\\\ &I&19 Dec 00&41&65 \\\\ &I&22 Dec 00&41&65 \\\\ &I&18 Jan 01&41&65 \\\\ &I&21 Jan 01&36&65\\\\ \\hline  \\end{tabular}  \\end{flushleft} \\caption{Observations} \\label{tab-obs} \\end{table*} ",
        "conclusions": "We have measured the Hubble constant $H_0$ in NGC~564 at $cz \\approx 5800$ km~s$^{-1}$ and in NGC~7619 at $cz \\approx 3700$ km~s$^{-1}$ with deep I--band SBF distance measurements at the ESO Very Large Telescope (VLT). We obtain $H_0 = 70 \\pm 7 \\pm 6$ km~s$^{-1}$/Mpc for NGC~564 and $H_0 = 68 \\pm 6 \\pm 6$ for NGC~7619. The actual SBF sample used for the measurement of $H_0$ in the Hubble Space Telescope Key Project on the Extragalactic Distance Scale (Freedman et al. 2001) amounts to six galaxies.  When we combine the measurements from this work with our previous VLT I--band SBF distance measurement in IC~4296 (Mei et al. 2000), we obtain : $H_0$ = 68 $\\pm$ 5 $\\pm$ 6 km~s$^{-1}$/Mpc. When we add the Freedman et al. (2001) SBF sample and calibrate our sample with the Ferrarese et al. (2000a) calibration, we obtain $H_0 = 71 \\pm 4 \\pm 6$ km~s$^{-1}$/Mpc. If we combine only galaxies beyond 4000~km~s$^{-1}$ from the two samples, we obtain $H_0$ = 72 $\\pm$ 6 $\\pm$ 6 km~s$^{-1}$/Mpc).  Our results are comparable with other I--band and K--band SBF measurements. More importantly we have increased the number of galaxies in the I--band sample with redshifts larger than 3000~km s$^{-1}$, and thereby reduced the statistical error on I--band SBF $H_0$ measurements from the Freedman et al. (2001) value of 5~km~s$^{-1}$/Mpc to 4~km~s$^{-1}$/Mpc (their measurement is in fact: $H_0$ = 70 $\\pm$ 5 $\\pm$ 6 km~s$^{-1}$/Mpc).  Few galaxies beyond the Fornax cluster have I--band SBF measurements and only five (included NGC~564 from this paper) have distances beyond 4000~km~s$^{-1}$. This lack of distant galaxies is responsible for the fact that the statistical error on $H_0$ obtained with SBF measurements is still comparable to the systematic uncertainty due to the Cepheid calibration, which is not the case for SNIa and TF methods (with a statistical error of, respectively,  2~km~s$^{-1}$/Mpc and 3~km~s$^{-1}$/Mpc (Freedman et al. 2001)) SBF being one of the most precise distance indicators, it is important to further reduce the statistical uncertainties. To achieve this, additional SBF measurements beyond 4000~km~s$^{-1}$ are needed."
    },
    "0212/astro-ph0212218_arXiv.txt": {
        "abstract": "The gravitational waves generated in the coalescence of massive binary black holes will be measurable by LISA to enormous distances. Redshifts $z \\sim 10$ or larger (depending somewhat on the mass of the binary) can potentially be probed by such measurements, suggesting that binary coalescences can be made into cosmological tools.  We discuss two particularly interesting types of probes. First, by combining gravitational-wave measurements with information about the universe's cosmography, we can study the evolution of black hole masses and merger rates as a function of redshift, providing information about the growth of structures at high redshift and possibly constraining hierarchical merger scenarios.  Second, {\\it if} it is possible to associate an ``electromagnetic'' counterpart with a coalescence, it may be possible to measure both redshift and luminosity distance to an event with less than $\\sim 1\\%$ error.  Such a measurement would constitute an amazingly precise cosmological standard candle.  Unfortunately, gravitational lensing uncertainties will reduce the quality of this candle significantly.  Though not as amazing as might have been hoped, such a candle would nonetheless very usefully complement other distance-redshift probes, in particular providing a valuable check on systematic effects in such measurements. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/astro-ph0212504_arXiv.txt": {
        "abstract": "Using survey data, we have re-evaluated the correlation of flat spectrum radio sources with \\EGRET sources in the Northern sky. An analysis incorporating the radio and X-ray properties and the $\\gamma$-ray source localization is used to gauge the reliability of associations and to search for counterparts of previously unidentified \\EGRET sources. Above $\\;\\vert b \\vert$=10$^\\circ$, where the classification is complete, we find that 70\\% of the Northern \\EGRET sources have counterparts similar to the bright \\EGRET blazars. For several of these we identify known blazar counterparts more likely than the earlier proposed 3EG association; for $\\sim$20 we have new identifications.  Spectroscopic confirmation of these candidates is in progress and we have found flat spectrum radio quasars and BL Lac counterparts with redshifts as high as 4. We also find strong evidence for a set of 28 objects with no plausible counterpart like the known \\EGRET blazars.  These thus represent either a new extragalactic population or a population of Galactic objects with a large scale height. The survey has been extended into the plane, where we find several new blazar candidates; the bulk of the sources are, however, Galactic. Looking ahead to the GLAST era, we predict that several of the present 3EG sources are composite and that higher resolution data will break these into multiple blazar IDs. ",
        "introduction": "The \\EGRET telescope on the {\\it CGRO} satellite has detected 271 sources in a survey of the Gamma-ray ($\\sim 100\\,$MeV to 10\\,GeV) sky. Of these roughly a quarter have been identified as blazars and along the Galactic plane there are a half dozen objects confirmed as young pulsars through their pulsed $\\gamma$-ray emission. Thus the bulk of the sources remain to be identified. There is a young Galactic population along the plane clearly correlated with high mass stars \\citep{kc96, yr97}. An intermediate latitude excess, especially in the direction of the Galactic bulge, suggests an older Galactic population whose nature is uncertain. Finally there remain a substantial number of high latitude sources with no AGN identification.  The selection of blazar candidates has largely proceeded by correlation with existing radio surveys \\citep{har99, mat01}. In contrast, the Galactic plane sources and individual intermediate latitude sources have been the subject of targeted multi-wavelength campaigns \\citep{rrk01, hal01, wal02}. In this project we attempt to obtain a more complete census of plausible blazar counterparts, sifting sources with extant radio survey data and then conducting a multi-wavelength follow-up. We are obtaining spectroscopic confirmation of the candidate AGN with Hobby$\\bullet$Eberly Telescope (HET) Marcario LRS spectroscopy. This survey has already discovered a number of new likely $\\gamma$-ray blazars, including good candidates for the most distant persistent $\\gamma$-ray sources known.  \\subsection{Blazar Properties}  The `blazar' label is somewhat heterogeneous, but in the context of the unified AGN model, these sources are believed to be viewed close to the axis of a powerful relativistic jet. As such they are compact flat spectrum radio sources, with apparent superluminal motion at VLBI scales. The optical counterparts exhibit significant polarization and OVV (optically violently variable) behavior \\citep{up95}. Optical spectroscopy often shows large equivalent width emission lines (flat spectrum radio quasars) while a significant fraction are continuum-dominated BL Lac-type objects. The broad-band spectral energy distribution (SED) is sometimes used to divide these into two classes with `red' blazars showing a synchrotron peak in the IR-optical with a Compton peak in the $\\gamma$-ray while `blue' blazars have a synchrotron component extending into the X-ray with a Compton peak inferred to extend to the TeV range (Urry 1999, and references therein). There are surveys underway to substantially increase the number of known blazars (e.g. DXRBS; Landt \\et \\,2001) which will eventually help in systematizing the broad-band properties of these sources.  Of these properties bright, flat spectrum radio emission seems best correlated with $\\gamma$-ray activity, but this may be primarily a selection effect of present counterpart lists, which focused on the brightest radio sources as the principal candidates.  In an effort to make a less biased census of the counterparts, we have re-examined the correlation with X-ray and radio properties in selecting candidates. ",
        "conclusions": "We believe that we have made a relatively unbiased assessment of the association of flat spectrum radio sources with \\EGRET AGN, identifying plausible counterparts down to 8.4GHz fluxes of $\\sim 100$mJy.  Assuming, as we have argued, that the bulk of our likely and plausible IDs are correct, we have substantially increased the completeness of identification of the Northern $|b|\\ge 10^\\circ$ \\EGRET sources from $\\sim$40\\% % to $\\sim 70$\\%. We have also argued that in a number of instances the 3EG sources are composites; including this, the increase in the number of proposed $\\gamma$-ray blazars is even larger.  Our identification allows selection of sources with substantially smaller 8.4GHz radio flux. It should not be too surprising that the number of candidates tapers off smoothly below the previous typical 1Jy value. We suspect that with smaller GLAST error ellipses or improved multi-wavelength constraints, identification could continue well below the $\\sim 0.1$Jy limit of our survey. Our identification of $\\sim 30$ sources that have flat spectrum counterparts (if any) well below 0.1Jy does not, of course, preclude that some of these may be radio-faint or steep spectrum AGN. Indeed, excess coincidences do continue slightly below our FoM identification limit (Figure 2). Also a number of the 3EG and \\citet{mat01} identifications not duplicated here are steep spectrum, but very bright low $z$ AGN. In these cases (with the exception of 3EG J0118+0248/3C037, discussed above) the steep spectrum proposed counterpart is located at large $\\Delta TS$. Nevertheless, if such sources are truly associated, it is reasonable that these form a distinct subset of the 3EG population, not identified in this analysis. However, the fact that an appreciable fraction of the `non-blazar' 3EG sources correlate with the Galactic disk suggests that many represent a new (old) Galactic population. Detailed assessment of the completeness of blazar IDs through the plane requires careful treatment of the 3EG exposure maps; we defer this and other population analysis to a later paper.  Along with fainter radio associations, we have pushed back the horizon of plausible blazar identifications (Figure 4). It will be interesting to compare the SED and VLBI $\\beta_\\perp$ properties of these fainter blazar counterpart candidates. This survey has found several individual targets of particular interest, including good $\\gamma$-ray source associations at $z=3-4$, several distant BL Lacs, and several low $z$ AGN that may be suitable targets for ground-based TeV observations. These targets should be the blazars brightest to GLAST and they will provide the most detailed spectra and light curves; as such they merit further study in preparation for the GLAST era. Since spectroscopic identifications are continuing, we defer detailed discussion of the red shift/ luminosity function distributions to a future paper.  The `non-blazars' are also excellent targets for further study in preparation for GLAST, since this sample will most likely produce new classes of high energy emitters.  While we do not discuss here the correlation of $\\gamma$-ray spectra and variability with the lower energy SEDs, we do note that our `non-blazars' correlate fairly well with the `steady' sources of \\citet{get00}  and the `persistent' sources of Grenier (2000), although exceptions occur. Clearly, pushing the improved identifications south to cover the Galactic bulge and beyond will be very important for characterization of these populations."
    },
    "0212/astro-ph0212397_arXiv.txt": {
        "abstract": "The properties of radio galaxies and quasars with and without optical or X-ray jets are compared.  The majority of jets from which high-frequency emission has been detected so far (13 with optical emission, 11 with X-rays, 13 with both) are associated with the most powerful radio sources at any given redshift.  It is found that optical/X-ray jet sources are more strongly beamed than the average population of extragalactic radio sources.  This suggests that the detection or non-detection of optical emission from jets has so far been dominated by surface brightness selection effects, not by jet physics.  It implies that optical jets are much more common than is currently appreciated. ",
        "introduction": "\\label{s:intro}  In the standard model for Active Galactic Nuclei (AGN), jets transfer mass, energy (kinetic and electromagnetic) and momentum from the central object into the surrounding medium.  \\citet{LiuZhang02} recently compiled a list of radio jets from the literature.  Adopting the definition from \\citet{BP84}, they list 925 jets associated with 661 radio sources.  Most of the jets have been identified by their radio synchrotron emission. Only a few of them also show optical synchrotron emission. Based on the most recent list of optical jets by \\citet{SU02} and the list of X-ray jets by \\citet{HK02}, I have compiled a list of 26 jets from 26 sources which have optical emission associated with them\\footnote{The full list with references is available online at \\url{http://home.fnal.gov/~jester/optjets} and will be contained in a forthcoming publication.  The following optical jets are not in \\citet{SU02}: 3C\\,31, PKS 0637$-$752, B2 0755+37, PKS 1136$-$135, PKS 1150+497, 3C\\,279, 3C\\,293, PKS 1354+195, 3C\\,303, B2 1553+254, B2 1658+30A, 3C\\,380.}.  Only a few optical jets have been studied in detail, most notably M87 \\citep[e.g.,][]{MRS96,Pereta99,Per01} and, in fewer bandpasses, 3C\\,273 \\citep[e.g.,][ and references therein]{RM91,JesterDiss}, but they have hardly been considered as a class of objects.  The small number of known optical jets suggests the conclusion that they are very rare objects.  This would, in fact, be expected from synchrotron physics: the magnetic fields expected in jets are of order of a few nanoTesla. This means that the loss timescales for electrons emitting optical synchrotron radiation are of order 1000 years or shorter, while typical arcsecond-scale jets have lengths from a few up to a few hundred kiloparsec.  Thus, it seems natural that the frequency $\\nu_\\mathrm{c}$ at which a jet's power-law spectrum cuts off lies below the optical wavelength region in most cases.  However, in most jets where optical emission has been observed, it extends over scales much larger than the lifetimes of ``optical'' electrons would permit.  \\citet{SU02} showed that \\emph{on average}, optical jets are sufficiently relativistic to remove the lifetime discrepancy by time dilation, without the need for reacceleration of particles within the jets \\citep[before this,][ had already shown that all optical jets known then are probably beamed]{Spaeta95}.  Still, at least in 3C\\,273's jet in-situ acceleration is necessary to explain the observed optical synchrotron emission, in particular because very strong beaming increases the energy loss by inverse Compton scattering of microwave background photons \\citep{Jes01}.  Moreover, the X-rays from many jets are explained as synchrotron emission as well \\citep{HK02} -- with much shorter loss timescales than in the optical, and correspondingly a more severe lifetime discrepancy which is harder to remedy by relativistic effects.  It appears thus that the number of optical and X-ray jets is small because some special physical conditions are necessary either to switch on the extended acceleration mechanism (whose nature remains elusive), or to accelerate the corresponding electrons near the AGN's core and allow their escape to kiloparsec scales.  However, in direct contradiction to this, \\citet{Capeta00} noted that the detection rate of optical jets in the B2 radio survey implies they are ``not particularly rare''.  Here I make an attempt to find out what the required physical conditions are, and indeed whether they are special.  Like \\citet{Spaeta95}, I find that beaming plays a role, and that more work is necessary before we can get beyond the observer-dependent beaming to the intrinsic physical conditions.\\footnote{The cosmology is that used by \\citet{LiuZhang02}: $H_0 = 100$\\,km/s/Mpc, $q_0 = 0.5$.} ",
        "conclusions": "\\label{s:conc}  The present comparison of the radio properties of sources with detected optical or X-ray jets to those with ``radio-only'' jets has not brought to light any connection between intrinsic physical properties and the presence of optical or X-ray emission.  Instead, it is found that optical jets are more strongly beamed on average than other radio jets are.  I therefore suggest that the entire sample of presently known optical jets is determined mainly by surface brightness selection effects.  Like \\citet{Pareta02} have done for the B2 sample, this suggestion has to be tested by comparing the optical surface brightness expected from an extrapolation of radio maps with a suitable power law ($f_\\nu \\propto \\nu^{-0.7}$ represents most jet spectra well) to the surface brightness limits of available optical observations.  For those jets without suitably deep observations, the most efficient search method will employ observations in the near-infrared. In this way, jets with cutoffs just below the optical are still detectable.  Positive detections will be followed up by observations in a second bandpass to determine the spectral index, and hence the detectability at higher frequencies.  Only with a firm sample of detections \\emph{and} non-detections can we analyse the dependence of the jet's synchrotron cutoff frequency $\\nu_\\mathrm{c}$ on other observables and thus address the question of what ignites optical jets.  \\hspace{4ex} I am grateful to Martin Hardcastle and Robert Laing for valuable comments.  \\setlength{\\bibsep}{0ex}"
    },
    "0212/astro-ph0212442_arXiv.txt": {
        "abstract": "The turbulent stresses that lead to angular momentum transport in accretion discs have often been treated as resulting from an isotropic effective viscosity, related to the pressure through the alpha parametrization of Shakura \\& Sunyaev.  This simple approach may be adequate for the simplest aspects of accretion disc theory, and was necessitated historically by an incomplete understanding of the origin of the turbulence.  More recently, Balbus \\& Hawley have shown that the magnetorotational instability provides a robust mechanism of generating turbulent Reynolds and Maxwell stresses in sufficiently ionized discs.  The alpha viscosity model fails to describe numerous aspects of this process.  This paper introduces a new analytical model that aims to represent more faithfully the dynamics of magnetorotational turbulent stresses and bridge the gap between analytical studies and numerical simulations.  Covariant evolutionary equations for the mean Reynolds and Maxwell tensors are presented, which correctly include the linear interaction with the mean flow.  Non-linear and dissipative effects, in the absence of an imposed magnetic flux and in the limit of large Reynolds number and magnetic Reynolds number, are modelled through five non-linear terms that represent known physical processes and are strongly constrained by symmetry properties and dimensional considerations.  The resulting model explains the development of statistically steady, anisotropic turbulent stresses in the shearing sheet, a local representation of a differentially rotating disc, in agreement with numerical simulations.  It also predicts that purely hydrodynamic turbulence is not sustained in a flow that adequately satisfies Rayleigh's stability criterion.  The model is usually formally hyperbolic and therefore `causal', and guarantees the realizability of the stress tensors.  It should be particularly useful in understanding the dynamics of warped, eccentric and tidally distorted discs, non-Keplerian accretion flows close to black holes, and a variety of time-dependent accretion phenomena. ",
        "introduction": "The magnetorotational instability is one of the most important instabilities in astrophysical fluid dynamics.  It applies to a differentially rotating, electrically conducting fluid in which the angular velocity decreases in magnitude away from the axis of rotation.\\footnote{In the absence of entropy gradients.  Otherwise, modified versions of the H\\o iland stability criteria exist (Papaloizou \\& Szuszkiewicz 1992; Balbus 1995).}  In the presence of a weak magnetic field of arbitrary configuration, such a flow is subject to a dynamical instability, with a growth rate comparable to the shear rate of the flow (Velikhov 1959; Chandrasekhar 1960; Fricke 1969; Acheson 1978; Balbus \\& Hawley 1991, 1992; Papaloizou \\& Szuszkiewicz 1992; Balbus 1995; Foglizzo \\& Tagger 1995; Ogilvie \\& Pringle 1996; Terquem \\& Papaloizou 1996).  The principal application of the magnetorotational instability is to accretion discs, in which the profile of angular velocity is fixed by Kepler's third law.  The non-linear development of the instability leads to sustained magnetohydrodynamic (MHD) turbulence, which transports angular momentum outwards in a vain attempt to neutralize the destabilizing gradient of angular velocity (Hawley, Gammie \\& Balbus 1995; Brandenburg et al. 1995; Stone et al. 1996; Balbus \\& Hawley 1998).  Owing to the generality of the conditions for instability, however, further applications exist to stellar interiors and other astrophysical objects, and possibly to geophysical situations.  Laboratory experiments are also in progress (Ji, Goodman \\& Kageyama 2001).  In one sense, the magnetorotational instability is generally considered to have `solved' the problem of the effective viscosity of accretion discs, through providing a robust means of angular momentum transport.\\footnote{Nevertheless, considerable interest and debate remain regarding the operation of the magnetorotational instability, or alternative mechanisms of angular momentum transport, in weakly ionized discs, e.g. around young stars and in the quiescent state of dwarf novae.}  However, the realization of the importance and ubiquity of the instability has, in some ways, made matters more difficult.  A strict view might be taken that any calculation of an accretion disc must now involve a numerical simulation of three-dimensional MHD turbulence.  These simulations are demanding to execute, complicated to analyse and still restricted in their scope and physical content.  To illustrate the limitations of the direct numerical approach, let us consider how the computational requirements scale with the parameters of the problem being studied.  Most simulations conducted so far have been based on the shearing box (Hawley et al. 1995), which is a local representation of a differentially rotating disc.  The simulated volume is comparable to $H^3$, where $H$ is the vertical scale-height or semi-thickness of the disc, and the models can be run for many times the dynamical (orbital) time-scale.  Let us accept, for the purposes of this argument, that recent versions of such simulations (e.g. Miller \\& Stone 2000) have achieved an adequate fidelity to physical reality (although questions indeed remain concerning numerical resolution and convergence, vertical boundary conditions, thermal and radiative physics, etc.).  How much more difficult would it be to simulate a global model of a realistic thin disc over many times the viscous time-scale, as would be required in a typical application?  Suppose that the disc has a constant angular semi-thickness $H/r$ and is simulated in a conical wedge in spherical polar coordinates $(r,\\theta,\\phi)$ with logarithmic grid spacing in $r$ between $r_{\\rm in}$ and $r_{\\rm out}$, and uniform spacing in $\\theta$ and $\\phi$.  To achieve a resolution comparable to that of a shearing box, the number of grid zones must increase by a factor of order $(r/H)\\ln(r_{\\rm out}/r_{\\rm in})$ in $r$, and a factor of order $r/H$ in $\\phi$.  The global viscous time-scale exceeds the dynamical time-scale in the inner part of the disc (where the Courant condition on the time-step is most severe) by a factor of order $(1/\\alpha)(r/H)^2(r_{\\rm out}/r_{\\rm in})^{3/2}$, where $\\alpha$ is the viscosity parameter of Shakura \\& Sunyaev (1973).  Therefore the computational requirements increase by a factor of order \\begin{equation} {{1}\\over{\\alpha}}\\left({{r}\\over{H}}\\right)^4 \\left({{r_{\\rm out}}\\over{r_{\\rm in}}}\\right)^{3/2} \\ln\\left({{r_{\\rm out}}\\over{r_{\\rm in}}}\\right). \\end{equation} Reasonable estimates of this factor for discs around young stars, in cataclysmic variables, and in X-ray binaries, respectively, are $10^{13}$, $10^{11}$ and $10^{15}$, or even larger.  (These numbers are reduced if the inner part of the disc is truncated by a stellar magnetosphere.)  For this reason, although some aspects of large-scale disc dynamics can now be addressed in direct numerical simulations (e.g.  Hawley 2001), global studies of realistic thin discs are currently inaccessible.  How then are we to study situations such as the thermal-viscous instability of discs in cataclysmic variables, the dynamics of warped or eccentric discs, tidally distorted discs in binary stars and protoplanetary systems, non-Keplerian accretion flows close to black holes, or the propagation of waves in discs?  The traditional approach, of course, has been to treat the turbulence as an effective viscosity given by the alpha parametrization of Shakura \\& Sunyaev (1973).  Despite suspicions that this treatment is likely to be inadequate in various respects, very few studies have aimed specifically at testing this issue (Abramowicz, Brandenburg \\& Lasota 1996; Torkelsson et al. 2000) and it is still unclear in what respects magnetorotational turbulence does or does not behave as an effective viscosity.  In a standard accretion disc the rate-of-strain tensor has a dominant $r\\phi$-component and an isotropic viscous model predicts that the $r\\phi$-component of the turbulent stress would similarly dominate.  This is not true of magnetorotational turbulence, in which other stress components (e.g. $\\phi\\phi$) are larger.  However, for the simplest problems such as the evolution of the surface density of a standard accretion disc, these additional stress components play essentially no role in the dynamics.  The alpha model is probably perfectly adequate in these circumstances, where one dimensionless number, $\\alpha$, is parametrizing a single physical quantity, the vertically integrated stress coefficient $T_{r\\phi}$.\\footnote{It may not be adequate for predicting the vertical distribution of energy dissipation, and therefore the vertical structure of the disc.}  In situations such as warped, eccentric, tidally distorted and non-Keplerian discs there is a large-scale, possibly slowly varying velocity field that is different from that in a standard disc.  From the perspective of a local observer, however, the flow resembles a shearing-box model, but with a non-standard (and possibly slowly varying) velocity gradient tensor.  Stress components other than the $r\\phi$-component can play a significant role in the dynamics of these systems, but whether the full turbulent stress tensor resembles a viscous stress in such situations is not understood.  In principle, shearing-box studies could be used to formulate a more faithful analytical or semi-analytical model of the relevant properties of the turbulence, which could then be used in studies of the large-scale behaviour of accretion discs.  This would bridge the gap between the local dynamics (on scales from the dissipation length up to $H$) and the large-scale dynamics, without requiring colossal advances in computation.  The formation and analysis of such a model is also likely to result in a better physical understanding.  In view of the formidable complexity of turbulent flows, such a programme might appear overambitious.  A complete theory of turbulence does not exist, because turbulent flows have an unlimited number of statistical properties that cannot be calculated from first principles.  However, in the present context, our attention is restricted to the question of how the mean turbulent stress in a patch of the disc responds to changes in the large-scale velocity gradient. As will be made clear, certain linear aspects of this problem can be deduced directly from the MHD equations; other non-linear aspects cannot be treated rigorously and require a closure model that is, nevertheless, strongly constrained by symmetry properties and dimensional considerations.  Although the model affords only an imperfect description of turbulence, its physically motivated nature and its connection to the MHD equations ensure that it performs more faithfully than the alpha viscosity model, and this is borne out by numerous test problems.  The remainder of this paper is organized as follows.  In Section~\\ref{A minimal model system for magnetorotational turbulence} I introduce a conceptual model system for magnetorotational turbulence, and make some simple arguments concerning the saturation of the turbulence.  I develop the equations for the mean Reynolds and Maxwell tensors as far as possible analytically in Section~\\ref{Stress equations and closures}, and then discuss the requirements of a non-linear closure model.  A specific model satisfying these principles is presented and explained physically in Section~\\ref{A simple model and its properties}.  The outcome of the model in the shearing sheet, under a variety of conditions, is investigated in Section~\\ref{Non-linear outcome in the shearing sheet}.  In Section~\\ref{The governing equations in compressible flows} the effects of a compressible mean flow are incorporated and the total energy budget is explained.  Some preliminary applications of the model are worked out in Section~\\ref{Applications and implications}. In Section~\\ref{Comparison with previous work} a comparison is made with existing models in engineering and in astrophysics, and also with published numerical simulations.  A summary follows in Section~\\ref{Conclusion}. ",
        "conclusions": "\\label{Conclusion}  The representation of the turbulent stresses in an accretion disc by the classical alpha viscosity model of Shakura \\& Sunyaev (1973) has been extremely successful in allowing the theory of accretion discs to develop even in the absence of a proper understanding of the turbulence, and in making possible a quantitative comparison between theory and observations.  However, the magnetorotational instability is now widely accepted to be the origin of the turbulence in sufficiently ionized discs (e.g. Balbus \\& Hawley 1998), and the alpha viscosity model fails to describe numerous aspects of this process. In this paper I have introduced a new analytical model that aims to represent more faithfully the dynamics of magnetorotational turbulent stresses and bridge the gap between analytical studies and numerical simulations.  Covariant evolutionary equations for the mean Reynolds and Maxwell tensors are presented, which correctly include the linear interaction with the mean flow.  Non-linear and dissipative effects, in the absence of an imposed magnetic flux and in the limit of large Reynolds number and magnetic Reynolds number, are modelled through five non-linear terms that represent known physical processes and are strongly constrained by symmetry properties and dimensional considerations.  The cooperation of these linear and non-linear effects explains, without any fine tuning, the development of statistically steady, anisotropic turbulent stresses in the shearing sheet, a local representation of a differentially rotating disc, in agreement with numerical simulations.  It reproduces other results seen in numerical studies: that purely hydrodynamic turbulence is not sustained in a flow that adequately satisfies Rayleigh's stability criterion; that the shear stress in magnetorotational turbulence in non-Keplerian discs is non-linearly related to the shear rate, at fixed angular velocity; and that the shearing epicyclic motions found in warped discs are damped by the turbulence at a slower rate than predicted on the basis of an isotropic viscosity.  It also makes predictions for future simulations, such as the decay of the two-dimensional wave and the possibility of sustained turbulence in situations where the angular velocity increases outwards.  The model has good formal properties, being typically hyperbolic and therefore `causal', guaranteeing the realizability of the stress tensors, and accounting satisfactorily for energy conservation.  Only the linear part of the model is derived directly from the MHD equations, and therefore the predictions of the model cannot be regarded as rigorous deductions.  In particular, the non-linear hydrodynamic stability of circular Keplerian flow remains unproven. Nevertheless, even without an accurate calibration of its parameters, the model almost certainly performs more faithfully than the alpha viscosity model in a variety of circumstances, and therefore represents an advance over previous work.  The general idea of seeking a model that connects the mean turbulent stress to changes in the mean velocity gradient can be related to one of the classical problems of continuum mechanics, which is to find the constitutive relation that characterizes a particular substance and allows the stress in a body to be related to its deformation history. From this perspective, the present model is reminiscent of certain classes of viscoelastic fluids.  In fact, the coexistence of the Reynolds and Maxwell stresses, which are influenced by the mean flow in different ways, suggests a kind of composite viscoelastic material. The reason that a constitutive relation might be thought to exist at all for magnetorotational turbulence is that the physics of the magnetorotational instability is local, suggesting that the mean stress in a volume comparable to $H^3$ should indeed depend only on the recent deformation history of that parcel of fluid.  In recent years, numerical simulations of accretion flows subject to the magnetorotational instability have been taken beyond the confines of the shearing box to study the effects of cylindrical geometry, the flow across the marginally stable orbit around a black hole, or the boundary layer between a star and a disc (e.g. Armitage 1998; Hawley 2000; Hawley, Balbus \\& Stone 2001; Armitage 2002).  These advances are to be welcomed, and represent significant computational achievements.  However, I believe there is still, and will continue to exist, a need for a description such as the one developed in this paper.  This is not just because global studies of realistic thin discs are well beyond the current capabilities of direct numerical simulations, but also because such a model will allow a variety of methods, analytical and less direct numerical ones, to continue to be used in studying accretion disc dynamics while taking seriously the nature of magnetorotational turbulence.  It should be particularly useful in understanding the dynamics of warped, eccentric and tidally distorted discs, non-Keplerian accretion flows close to black holes, and a variety of time-dependent accretion phenomena."
    },
    "0212/astro-ph0212456_arXiv.txt": {
        "abstract": "The evolution of magnetic fields frozen to a perfectly conducting plasma fluid around a Kerr black hole is investigated. We focus on the plunging region between the black hole horizon and the marginally stable circular orbit in the equatorial plane, where the centrifugal force is unable to stably balance the gravitational force. Adopting the kinematic approximation where the dynamical effects of magnetic fields on the fluid motion are ignored, we exactly solve Maxwell's equations with the assumptions that the geodesic motion of the fluid is stationary and axisymmetric, the magnetic field has only radial and azimuthal components and depends only on time and radial coordinates. We show that the stationary state of the magnetic field in the plunging region is uniquely determined by the boundary conditions at the marginally stable circular orbit. If the magnetic field at the marginally stable circular orbit is in a stationary state, the magnetic field in the plunging region will quickly settle into a stationary state if it is not so initially, in a time determined by the dynamical time scale in the plunging region. The radial component of the magnetic field at the marginally stable circular orbit is more important than the toroidal component in determining the structure and evolution of the magnetic field in the plunging region. Even if at the marginally stable circular orbit the toroidal component is zero, in the plunging region a toroidal component is quickly generated from the radial component by the shear motion of the fluid. Finally, we discuss the dynamical effects of magnetic fields on the motion of the fluid in the plunging region. We show that the dynamical effects of magnetic fields are unimportant in the plunging region if they are negligible on the marginally stable circular orbit. This supports the ``no-torque inner boundary condition'' of thin disks, contrary to the claim in the recent literature. ",
        "introduction": "\\label{sec1}  It is widely believed that accretion disks around Kerr black holes exist in many astrophysical environments, ranging from active galactic nuclei to some stellar binary systems \\cite{fra02,kat98}. People usually assume that the inner boundary of a thin Keplerian disk around a Kerr black hole is located at the marginally stable circular orbit, inside which the centrifugal force is unable to stably balance the gravity of the central black hole \\cite{nov73,pag74}. In the disk region, particles presumably move on nearly circular orbits with a very small inward radial velocity superposed on the circular motion, the gravity of the central black hole is approximately balanced by the centrifugal force. As disk particles reach the marginally stable circular orbit, the gravity of the central black hole becomes to dominate over the centrifugal force and the particles begin to nearly free-fall inwardly. The motion of fluid particles in the plunging region quickly becomes supersonic then the particles loose causal contact with the disk, as a result the torque at the inner boundary of the disk is approximately zero \\cite[and references therein]{muc82,abr89,pac00,arm01,afs02}. This is usually called the ``no-torque inner boundary condition'' of thin accretion disks.  Some recent studies on the magnetohydrodynamics (MHD) of accretion disks have challenged the ``no-torque inner boundary condition''. Magnetic fields have been demonstrated to be the most favorable agent for the viscous torque in an accretion disk transporting angular momentum outward \\cite[and references therein]{bal98}. By considering the evolution of magnetic fields in the plunging region, Krolik \\cite{kro99} pointed out that magnetic fields can become dynamically important in the plunging region even though they are not so on the marginally stable circular orbit, and argued that the plunging material might exert a torque to the disk at the marginally stable circular orbit. With a simplified model, Gammie \\cite{gam99} solved Maxwell's equations in the plunging region and estimated the torque on the marginally stable circular orbit. He demonstrated that the torque can be quite large and thus the radiation efficiency of the disk can be significantly larger than that for a standard accretion disk where the torque at the inner boundary is zero. Furthermore, Agol and Krolik \\cite{ago00} have investigated how a non-zero torque at the inner boundary affects the radiation efficiency of a disk.  Numerical simulations of MHD disks \\cite{haw00,arm01,haw01a,haw01b,haw02} have greatly improved our understanding of disk accretion processes. These simulations show that the magneto-rotational instability effectively operates inside the disk and leads to accretion, though the accretion picture is much more complicated than that assumed in the standard theory of accretion disks. Generally, the disk accretion is non-axisymmetric and strongly time-dependent. It is also found that, as disk material gets into the plunging region, the magnetic stress at the marginally stable circular orbit does not vanish but smoothly extends into the plunging region \\cite{haw01a,haw01b,haw02}, though the effect is significantly reduced as the thickness of the disk goes down \\cite{arm01,haw01a}. Furthermore, the specific angular momentum of particles in the plunging region does not remain constant, which implies that the magnetic field may be dynamically important in the plunging region \\cite{haw01a,haw01b,haw02}. All these results are fascinating and encouraging. Unfortunately, due to the limitation in space resolution and time integration, stationary and geometrically thin accretion disks are not accessible to the current 2-D and 3-D simulations. So it remains unclear how much insights we can get for stationary and geometrically thin accretion disks from these simulations \\cite{afs02}.  Instead of small-scale and tangled magnetic fields in an accretion disk transporting angular momentum within the disk, a large-scale and ordered magnetic field connecting a black hole to its disk may also exist and play important roles in transportation of angular momentum and energy between the black hole and the disk \\cite{bla98,li00a,li00b,li02a,li02b}. Recent {\\it XMM-Newton} observations of some Seyfert galaxies and Galactic black hole candidates provide possible evidences for such a magnetic connection between a black hole and its disk \\cite{wil01,li01,mil02}.  All these promote the importance of studying the evolution and the dynamical effects of magnetic fields around a Kerr black hole.  In this paper, we use a simple model to study the evolution of magnetic fields in the plunging region around a Kerr black hole. We assume that around the black hole the spacetime metric is given by the Kerr metric; in the plunging region, which starts at the marginally stable circular orbit and ends at the horizon of the black hole, a stationary and axisymmetric plasma fluid flows inward along timelike geodesics in a small neighborhood of the equatorial plane. The plasma is perfectly conducting and a weak magnetic field is frozen to the plasma. The magnetic field and the velocity field have two components: radial and azimuthal. We will solve the two-dimensional Maxwell's equations where the magnetic field depends on two variables: time and radius, and investigate the evolution of the magnetic field. This model is similar to that studied by Gammie \\cite{gam99}, but here we include the time variable. Furthermore, we ignore the back-reaction of the magnetic field on the motion of the plasma fluid to make the model self-consistent, since if the dynamical effects of the magnetic field are important the strong electromagnetic force will make the fluid expand in the vertical direction. The ignorance of the back-reaction of the magnetic field will allow us to to analytically study the evolution of the magnetic field, but it will prevent us from quantitatively studying the dynamical effects of the magnetic field. However, we believe that the essential features of the evolution of the magnetic field in the plunging region is not sensitive to the details of the dynamical effects. To check the self-consistency of the model, we will estimate the dynamical effects of the magnetic field by considering the back-reaction of the magnetic field on the fluid motion. We will also discuss the self-consistent solutions to the coupled Maxwell and dynamical equations, and look for implications to the ``no-torque inner boundary condition''.  The paper is organized as follows: In Sec. \\ref{sec2} we write down the Maxwell's equations for an ideal MHD plasma. In Sec. \\ref{sec3} we write down the general forms of the magnetic field and the velocity field of the plasma fluid around a Kerr black hole. In Sec. \\ref{sec4} we solve Maxwell's equations with the approximations outlined above. In Sec. \\ref{sec5}, with the solutions obtained in Sec. \\ref{sec4}, we study the evolution of the magnetic field in the plunging region. In Sec. \\ref{sec6}, we check if magnetic fields can become dynamically important in the plunging region. In Sec. \\ref{sec7} we draw our conclusions. In the Appendix we solve Maxwell's equations in the disk region, since we want our solutions in the plunging region to be continuously joined to the solutions in the disk region.  Throughout the paper we use the geometrized units $G = c = 1$ and the Boyer-Lindquist coordinates $(t,r,\\theta,\\phi)$ \\cite{mis73,wal84}, except in the Appendix where cylindrical coordinates are used. ",
        "conclusions": "\\label{sec7}  With a simple model we have studied the evolution of magnetic fields in the plunging region around a Kerr black hole. The model contains the following assumptions: (1) The background spacetime is described by the Kerr metric; (2) The plasma is perfectly conducting so that the magnetic field is frozen to the plasma fluid; (3) The kinematic approximation \\cite{bal98} applies, i.e. the dynamical effects of the magnetic field on the fluid motion are negligible so that the plasma fluid flows along timelike geodesics toward the central black hole; (4) In a small neighborhood of the equatorial plane (i.e., $|\\pi/2 - \\theta| \\ll 1$) the magnetic field and the velocity field have only radial and azimuthal components. The assumption (4) implies that $u^\\theta = B^\\theta = 0$ and $\\partial u^\\theta/\\partial \\theta = \\partial B^\\theta /\\partial \\theta = 0$ on the equatorial plane. With above assumptions, we have exactly solved Maxwell's equations for axisymmetric solutions (i.e. $\\partial/\\partial \\phi = 0$). The solutions are given by Eqs.~(\\ref{fsol1}) and (\\ref{fsol2}) [or, equivalently, Eqs.~(\\ref{br}) and (\\ref{bf})], where $C_0$ is a constant measuring the magnetic flux through a circle in the equatorial plane, $\\Psi$ is a function determining the time-evolution of the magnetic field. Both $C_0$ and $\\Psi$ are determined by the boundary conditions.  The dependence of the solutions on the coordinate time $t$ is given by the function $\\Psi = \\Psi (t-\\tau)$, where $\\tau = \\tau(r)$ is defined by Eq.~(\\ref{tau}). The function $\\Psi$ is uniquely determined by the initial and boundary conditions of the magnetic field. The general form of $\\Psi$ determines the retarded nature of the solutions: at any time the state of the magnetic field at a radius in the plunging region is determined by the state of the magnetic field at a larger radius and an earlier time (Figs.~\\ref{fig1} and \\ref{fig2}). This suggests that the stationary state of the magnetic field in the plunging region is uniquely determined by the boundary conditions at the marginally stable circular orbit.  Examples for the evolution of magnetic fields in the plunging region are shown in Figs.~\\ref{fig3} -- \\ref{fig8}, for both the stationary case ($\\Psi = {\\rm constant}$, Figs.~\\ref{fig3} -- \\ref{fig5}) and the non-stationary case ($\\Psi \\neq {\\rm constant}$, Figs.~\\ref{fig6} -- \\ref{fig8}). From these figures we see that, the boundary value of the radial component of the magnetic field at the marginally stable circular orbit is more important in determining the strength of the magnetic field in the plunging region than the toroidal component. The initially toroidal component is attenuated in the plunging region by the radial expansion of the fluid (Fig.~\\ref{fig3}). The initially radial component is amplified by the azimuthal and vertical compression (the convergence of the fluid), and a toroidal component is generated from the radial component by the shear motion of the fluid though the toroidal component is initially zero at the marginally stable circular orbit (Fig.~\\ref{fig4}). This leads to the amplification of the magnetic field in the plunging region (Fig.~\\ref{fig5}). The evolution of the magnetic field depends on the spinning state of the black hole, but is more sensitive to the initial value of the radial velocity on the marginally stable circular orbit [or, equivalently the parameter $\\delta$ defined by Eq.~(\\ref{le0})].  The time-evolution of magnetic fields shown in Figs.~\\ref{fig6} -- \\ref{fig8} confirms our claim that the stationary state of the magnetic field in the plunging region is uniquely determined by the boundary conditions on the marginally stable circular orbit. If in the plunging region a magnetic field is initially deviated from the stationary solutions, it will evolve to and finally get saturated at the state given by the stationary solutions in a dynamical time scale determined by the free-fall motion of fluid particles in the plunging region. If the magnetic field on the marginally stable circular orbit is in a stationary state, the magnetic field in the plunging region will automatically settle into a stationary state. Thus, the evolution of magnetic fields in the plunging region is very different from that in the disk region, in the latter case the state of the magnetic field is determined by the local MHD processes. The difference is caused by the following fact: in the disk region the fluid has a very small radial velocity so that local MHD processes have shorter time scales than the radial motion; in the plunging region the fluid has a large radial velocity so that local MHD processes usually have longer time scales than the radial motion.  In deriving the solutions we have assumed that the dynamical effects of the magnetic field in the plunging region are unimportant [the kinematic approximation adopted in assumption (3)]. To justify this assumption, we have studied the dynamical effects of magnetic fields in the stationary state in two ways. First, we estimate the dynamical effects of magnetic fields by considering the back-reaction: substituting the solutions of Maxwell's equations we obtained, where we assumed that fluid particles move on geodesics, into the dynamical equations to check if the motion of the fluid is significantly affected by magnetic fields. In this way, we estimate the dynamical effects of magnetic fields on the motion of the fluid by calculating the parameters $\\eta_1$, $\\eta_2$ and $\\eta_3$ defined by Eq.~(\\ref{eta}). The mass density of the fluid, which appears in the denominators of $\\eta_1$, $\\eta_2$ and $\\eta_3$, drops quickly in the plunging region if $|v_{\\hat{r}}| \\ll |v_{\\hat{\\phi}}|$ on the marginally stable circular orbit (Fig.~\\ref{fig9}). However, the evolution of the magnetic field, which appears in the numerators of $\\eta_1$, $\\eta_2$ and $\\eta_3$, sensitively depends on the orientation of the magnetic field on the marginally stable circular orbit, the latter is essentially determined by the MHD processes in the disk region. If we require that the solutions in the plunging region are smoothly joined to the solutions in the thin Keplerian disk region, then it is reasonable to assume that on the marginally stable circular orbit, as well as in the disk region, the magnetic field is parallel to the velocity field: $B^r/B^\\phi = u^r/u^\\phi$ (see the Appendix). This implies that for the solutions in the plunging region we should have $\\Psi = 0$ [Eq.~(\\ref{psi0})], i.e. the orientation of the magnetic field follows the orientation of the velocity field of the fluid [Eq.~(\\ref{pc})]. With such a boundary condition, we have calculated $\\eta_1$, $\\eta_2$ and $\\eta_3$\\,, assuming that particles move on geodesics. We see that, if they are $\\ll 1$ on the marginally stable circular orbit, then they remain so in the plunging region (Figs.~\\ref{fig10} and \\ref{fig11}). Indeed, $\\eta_1$, $\\eta_2$ and $\\eta_3$ decrease in the plunging region for sufficiently small $|v_{\\hat{r}}|$ on the marginally stable circular orbit.  Second, we self-consistently solve the coupled Maxwell and dynamical equations on the horizon of the black hole, check how much has changed in the specific angular momentum of fluid particles since they leave the marginally stable circular orbit. The specific energy of fluid particles is not changed by a magnetic field that satisfies $\\Psi = 0$ [see Eq.~(\\ref{fe})]. However, the specific angular momentum does change [see Eq.~(\\ref{fl})]. We have calculated the specific angular momentum of fluid particles as they reach the horizon of the black hole, assuming that the fluid particles pass the fast critical point near the marginally stable circular orbit. The deviation from the specific angular momentum as the particles just leave the marginally stable circular orbit is shown in Fig.~\\ref{fig12}. We see that, though the specific angular momentum of fluid particles is changed by the magnetic field, the effects are always small assuming that on the marginally stable circular orbit the dynamical effects of the magnetic field are not important. For example, for $a/M \\ge 0$ we have $|\\Delta L/L| < 3\\%$ if $\\eta_1 <10^{-2}$ on the marginally stable circular orbit. We have also calculated the ratios $\\eta_1 = B^2/4\\pi\\rho_{\\rm m}$ and $c_{{\\rm A}r} c_{\\rm A}^{~\\,r}/u_r u^r$ on the horizon (Fig.~\\ref{fig13}), where $c_{\\rm A}^{~\\,r}$ is the radial component of the Alfv\\'{e}n velocity. These two ratios are relevant to the critical points in the flow [see Eqs.~(\\ref{crur}), (\\ref{car}), and (\\ref{eta1})] and measure at what level the magnetic field affects the motion of the fluid. Fig.~\\ref{fig13} shows that their values on the horizon are quite small.  Both approaches (back-reaction and self-consistent solutions) confirm that the dynamical effects of the magnetic field are unimportant in the plunging region if they are so on the marginally stable circular orbit.  Our results differ from that of Gammie \\cite{gam99}, in which he claimed that in the plunging region the dynamical effects of magnetic fields can be important. The difference is caused by the fact that Gammie used a boundary condition that is different from ours for the magnetic field. Gammie assumed that $\\Psi = - \\Omega_{\\rm ms} C_0$, where $\\Omega_{\\rm ms}$ is the angular velocity of the disk at the marginally stable circular orbit, while we assume that $\\Psi = 0$. Gammie's boundary condition implies that $v_{\\hat{r}} = 0$ at $r = r_{\\rm ms}$ (as clearly stated in his paper), which makes $r = r_{\\rm ms}$ a singular point where $\\rho_{\\rm m} = \\infty$ (to keep the mass flux $F_{\\rm m}$ nonzero). Hence, Gammie's solutions are not well-behaved at the marginally stable circular orbit. Our boundary condition allows the solutions in the plunging region to be smoothly joined to the solutions in the disk region (see the Appendix), since in our solutions all physical quantities are finite at the marginally stable orbit. Certainly, to precisely take care of the transition from the disk region to the plunging region, gas pressure and non-electromagnetic stress must be properly taken into account near the inner edge of the disk.  As mentioned in the Introduction, recently the ``no-torque inner boundary condition'' for thin accretion disks has been challenged by some authors (including Gammie) based on their studies on the evolution of magnetic fields in the plunging region \\cite{kro99,gam99,haw01b,haw02}. The results in this paper suggest that in the plunging region the dynamical effects of the magnetic field are not important, if the solutions in the plunging region are smoothly joined to the the solutions in the thin Keplerian disk region. Thus, the argument against the ``no-torque inner boundary condition'' is not founded.  Finally, we note that in the paper we have neglected dissipative processes like magnetic reconnection and ohmic dissipation, and the evaporation effect arising from the magnetic buoyancy and MHD instabilities. These processes operate in disks to limit the amplification of magnetic fields \\cite[and references therein]{gal79,tou92,bal98}. Certainly it is conceivable that they can also operate in the plunging region to reduce the amplification effect of magnetic fields. If these processes are important, the results of this paper tend to overestimate the amplification of magnetic fields in the plunging region. Then, the dynamical effects of magnetic fields should be weaker than that we have estimated without considering dissipative processes, which will strengthen our conclusion that the dynamical effects of magnetic fields are unimportant in the plunging region."
    },
    "0212/hep-ph0212213_arXiv.txt": {
        "abstract": "In non-oscillatory (NO) inflationary models, the reheating mechanism was usually based on gravitational particle production or the mechanism of instant preheating. In this paper we introduce the curvaton mechanism into NO models to reheat the universe and generate the curvature perturbations. Specifically we consider the Peebles-Vilenkin quintessential inflation model, where the reheating temperature can be extended from $1$MeV to $10^{13}$GeV. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/astro-ph0212089_arXiv.txt": {
        "abstract": "{Using full general relativistic calculations, we investigate the possibility of generation of mass outflow from spherical accretion onto non-rotating black holes. Introducing a relativistic hadronic-pressure-supported steady, standing, spherically-symmetric shock surface around a Schwarzschild black hole as the effective physical barrier that may be responsible for the generation of spherical wind, we calculate the mass outflow rate $R_{\\dot m}$ in terms of three accretion parameters and one outflow parameter by simultaneously solving the set of general relativistic hydrodynamic equations describing spherically symmetric, transonic, polytropic accretion and wind around a Schwarzschild black hole. Not only do we provide a sufficiently plausible estimation of $R_{\\dot m}$, we also successfully study the dependence and variation of this rate on various physical parameters governing the flow. Our calculation indicates that independent of initial boundary conditions, the baryonic matter content of this shock-generated wind always correlates with post-shock flow temperature. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/astro-ph0212040_arXiv.txt": {
        "abstract": "The hydrodynamics and intrinsic properties of galactic-scale gaseous outflows generated by violent starbursts are thoroughly discussed, taking into account the hot gas chemical evolution and  radiative cooling. We also discuss the observational properties of supergalactic winds in X-rays and visible line regimes, derived from the hydrodynamic calculations. ",
        "introduction": "Massive starbursts, well localized short episodes of violent star formation, are among the intrinsic characteristics of the evolution of galaxies. They are found both at high and intermediate  redshifts as well as in galaxies in the Local universe.  The large kinetic luminosity produced by violent bursts of star formation is well known to drastically affect the surrounding interstellar medium (ISM), generating giant superbubbles, supershells and in extreme cases supergalactic winds. Powerful gas outflows from star forming regions are important redistributors of the ISM mass and momentum. On the other hand, the gas ejected by supernova explosions contains the products of massive stars. Therefore gaseous outflows initiated by starburst affect also the chemical evolution of the galactic ISM and the intergalactic medium.  Supergalactic winds are expected to be detectable thanks to their extended X-ray emission. Indeed, diffuse X-ray emission associated with starburst galaxies has been detected by ROSAT, ASCA, BeppoSAX, Chandra and XMM missions in several starburst sources. However, the origin of the soft X-ray emission, its connection with H$_{\\alpha}$ filaments and the chemical composition of the X-ray emitting gas, remain ambiguous.  The paper deals with three central aspects of the resultant outflows. The chemical evolution expected in the interior of superbubbles. The conditions required to establish a supergalactic wind and the intrinsic properties and observational manifestations of fully developed supergalactic winds. ",
        "conclusions": "The results from the calculations presented here imply that:  The metallicity of superbubbles vary with time and can easily exceed the solar value even if the host galaxy ISM has a low metal abundance. The effects of metal contamination are most noticeable during the first 10 Myr of the bubble evolution.  Galactic superwinds driven by compact and powerful starbursts undergo catastrophic cooling close to the star cluster surface and establish a temperature distribution radically different to that predicted by adiabatic calculations.  The fall of the superwind temperature leads to a smaller zone radiating in X-rays and decreases the superwind X-ray luminosity.  At the same time cooling brings the inner radius of the warm ionized gas envelope closer to the star cluster. This increases the estimated H$_{\\alpha}$ luminosity and predicts a low-intensity broad ($\\sim 1000$ km s$^{-1}$) line emission component."
    },
    "0212/astro-ph0212276_arXiv.txt": {
        "abstract": "We present constraints on several cosmological parameters from a combined analysis of the most recent Cosmic Microwave Background anisotropy data and the Sloan Digital Sky Survey cluster mass function. We find that the combination of the two data sets breaks several degeneracies among the parameters and provides the following constraints: $\\sigma_8=0.76\\pm0.09$, $\\Omega_m=0.26^{+0.06}_{-0.07}$, $h=0.66^{+0.05}_{-0.06}$, $n=0.96 \\pm 0.05$, $\\tau_c=0.07^{+0.07}_{-0.05}$. ",
        "introduction": "The last few years have seen a spectacular increase in the amount and quality of available cosmological data. The new results on the Cosmic Microwave Background angular power spectrum \\citep{netterfield, halverson, lee, pearson, scott, benoit} have confirmed the theoretical prediction of acoustic oscillations in the primeval plasma and constrained theories of large-scale structure formation \\citep[see e.g.][]{wang}. At the same time, early data from the 2dFGRS \\citep{2dfgrs} and the Sloan Digital Sky Survey \\citep{yor00, sto02} galaxy redshift surveys have provided an unprecedented view of the large-scale structure of the universe as traced by galaxies.  Combined analysis of these independent CMB and galaxy data sets are placing strong constraints on some of the fundamental cosmological parameters \\citep{bah99,efstathiou, lahav, mesilk}. Together with the high-redshift supernovae results \\citep{perlmutter, filippenko},  a concordance model of a flat, low-density Cold Dark Matter cosmology has become the current paradigm.  The goal of these analyses is to determine the precise values of the cosmological parameters of the $\\Lambda$-CDM model. Due to 'cosmic degeneracy', the CMB data alone are unable to place tight constraints on several fundamental parameters, such as the r.m.s. amplitude of the mass fluctuations $\\sigma_8$, the Hubble parameter $h$ and the optical depth of the universe $\\tau_c$, even if one restricts the analysis to a flat universe.   In the present {\\it Letter} we combine the most recent CMB anisotropies data with the constraints obtained from the mass function of clusters of galaxies determined from early commissioning imaging data of the SDSS \\citep{SDSSmf03} to break the degeneracy among the cosmological models, allowing a determination of the best-fit values for individual parameters. ",
        "conclusions": "We combine the Cosmic Microwave Background anisotropy and SDSS cluster mass function data to produce constraints on several cosmological parameters. The complementary nature of the two data sets breaks existing degeneracies among cosmological parameters. The CMB data tend to indicate a higher $\\Omega_m$ and a lower amplitude $\\sigma_8$ than suggested by the cluster data; thus the combined result for $\\Omega_m$ is pulled to a lower value than suggested by the CMB alone, and the amplitude $\\sigma_8$ is lower than suggested by the cluster data alone. We find the combined data suggest a mass density $\\Omega_m=0.26_{-0.07}^{+0.06}$; a normalization of the matter power spectrum $\\sigma_8=0.76\\pm0.09$; Hubble parameter $h=0.66^{+0.05}_{-0.06}$; a nearly scale-invariant spectrum of primordial fluctuations $n=0.96 \\pm 0.05$; and an optical depth of the universe $\\tau_c=0.07^{+0.07}_{-0.05}$. We have restricted the analysis to flat universes; if this is the case, the density of mass-energy in the universe is dominated by a form other than ordinary matter."
    },
    "0212/astro-ph0212106_arXiv.txt": {
        "abstract": "Fourier decomposition is a well established technique used in stellar pulsation. However the quality of reconstructed light curves using this method is reduced when the observed data have uneven phase coverage. We use simulated annealing techniques together with Fourier decomposition to improve the quality of the Fourier decomposition for many OGLE LMC fundamental mode Cepheids. This method restricts the range that Fourier amplitudes can take. The ranges are specified by well sampled Cepheids in the Galaxy and Magellanic Clouds.  We also apply this method to reconstructing Cepheid light curves observed by the HST. These typically consist of 12 V band and 4 I band points. We employ a direct Fourier fit to the 12 V band points using the simulated annealing method mentioned above and explicitly derive and use Fourier interrelations to reconstruct the I band light curve. We discuss advantages and drawbacks of this method when applied to HST Cepheid data over existing template methods. Application of this method to reconstruct the light curves of Cepheids observed in NGC 4258 shows that the derived Cepheid distance ($\\mu_0=29.38\\pm0.06$, random error) is consistent with its geometrical distance ($\\mu_0=29.28\\pm0.09$) derived from observations of water maser. ",
        "introduction": "The study of Cepheid light curves has two major applications: (a) distance determinations to nearby galaxies via the well calibrated period-luminosity (PL) relations (for example, see \\citet{mad85,fea87,mad91,sah00c}); (b) the understanding of the pulsational behavior of Cepheids by comparing the observed light curves to the theoretical counterparts \\citep{sim83,buc90}. Therefore, reconstructing the Cepheid light curves from the observed photometric data is important regarding these two aspects.  Because Cepheid light curves are periodic, the data points from well observed Cepheids can be described by Fourier expansion. This technique was first introduced by Schaltenbrand \\& Tammann in 1971, and developed by \\citet{sim81} to study the structural properties of Cepheid light curves. In general, the $n^{th}$-order Fourier expansion has the following form:  \\begin{eqnarray} m(t) = A_0 + \\sum_{j=1}^{n}A_j\\cos(j\\omega t + \\phi_j), \\end{eqnarray}  where $\\omega=2\\pi /P$ is the frequency and $P$ is the period of Cepheid. The $A_0$ term is the mean value of the light curve; the $A_j$ and $\\phi_j$ are the Fourier amplitudes and phases for $j^{th}$ order, respectively. In this paper we develop the use of simulated annealing techniques to improve the quality of light curve reconstruction in nearby galaxies such as the LMC. We also apply these techniques in the light curve reconstruction of Cepheids observed by Hubble Space Telescope (HST) which is necessary for distance determinations. In particular, we use Fourier expansion, as given by equation (1) but with modifications, to reconstruct the V band light curves and Fourier interrelations to reconstruct the I band light curves. These Fourier techniques have been established in stellar pulsation studies. We also present some potential advantages and drawbacks of the application of our method to this problem.  Section 2 \\& 4 will describe the Fourier expansion and interrelations in details, respectively. Section 3 briefly describe the application of these Fourier techniques to HST data. The error analysis of the light curve reconstruction procedures will be presented in Section 5, following by the discussion and conclusion in Section 6. In appendix, we show another Fourier method, the Fourier intrarelations, that can be applied to reconstruct the Cepheid light curves. However, this method is less preferable than Fourier interrelations. ",
        "conclusions": "The light curve of a Cepheid can be reconstructed with the Fourier techniques described in this paper. For OGLE LMC Cepheid data, we have demonstrated that we can frequently improve the quality of the reconstructed light curve by restricting the range of Fourier amplitudes can take in the Fourier fit. These ranges are obtained from well sampled Cepheid light curves in the Galaxy and Magellanic Clouds.  We have then applied these techniques to HST Cepheid data. Such data typically consist of 12 and 4 V and I band points respectively. In particular, Cepheids with sufficient observations permit a $4^{th}$ order Fourier expansion. However, the Cepheids in HST I band observations that have insufficient data points require the application of Fourier interrelations to reconstruct the I band light curves from the V band. In summary, the Fourier techniques for light curve reconstruction procedures, as applied to the HST data, include two main parts:  \\begin{itemize} \\item  $4^{th}$ order Fourier expansion to 12 V band data points. This is reconstructing the light curves by a direct fit to the data points, with appropriate ranges for each of the Fourier parameters. \\item  Fourier interrelation to 4 I band data points. This is using the linear relations of Fourier parameters between V and I bands to reconstruct the light curves. \\end{itemize}  In general, these relatively simple techniques can reconstruct the Cepheid light curves quite well, given that the data points are relatively uniform and well sampled. These techniques serve as an alternative method to obtain the mean magnitude of Cepheids besides phase weighted intensity mean and template fitting procedures. However, the detailed comparisons between these methods are beyond the scope of this paper. Some examples of the reconstructed light curves for extra-galactic Cepheids (observed by H0KP) with these Fourier techniques are presented in Figure \\ref{fig:fig19}, which show good agreements to the observed data points.  We also performed Monte-Carlo simulations to determine the errors associated with our light curves reconstruction procedures. The errors of the mean magnitudes and Fourier amplitudes are determined from the fits of Gaussian distribution to the error histograms. These Gaussian fits show that our procedure is unbiased. The errors are: \\begin{itemize} \\item\tV band: $\\sigma_{mean}=0.034$; $\\sigma_{A1}=0.039$; $\\sigma_{A2}=0.033$; $\\sigma_{A3}=0.028$ and $\\sigma_{A4}=0.023$. \\item\tI band: $\\sigma_{mean}=0.035$; $\\sigma_{A1}=0.025$; $\\sigma_{A2}=0.021$; $\\sigma_{A3}=0.017$ and $\\sigma_{A4}=0.014$. \\end{itemize}  Could it be the case that the distribution of 12 points is so severe that a decent fit to the V band data is not possible? Certainly and it could be that in this case a template type technique will produce a V and I band mean. However, we again contend that in this case, it will probably be the case that this point source is not used in the final analysis.  \\subsection{The Correlations of Fourier Amplitudes}  It has been suggested by anonymous referee that the higher order Fourier amplitudes shown in Figure \\ref{fig:fig3} \\& \\ref{fig:fig5} are correlated with the first order Fourier amplitude ($A_1$) in the same bands (the Fourier intrarelations). These have been examined, for example, by \\citep{ant87} and in the appendix. However, the Fourier intrarelations are less preferable than Fourier interrelations for reconstructing the Cepheid light curves. The reasons are discussed in following examples and in Appendix.  Here, we reexamine the results of \\citet{ant87} with our data sets and compare them to the Fourier interrelations. We used all ``calibrating set'' Cepheids and pick only the ``good'' OGLE LMC Cepheids (crosses) from Figure \\ref{fig:fig3}. Then we re-plot the $A_1-A_2$ Fourier intrarelations and the $A_1(V)-A_1(I)$ Fourier interrelation in Figure \\ref{fig:fig20}. The ``Calibrating set'' Cepheids and OGLE LMC Cepheids are in left and right panels, respectively. In the figure, we distinguished Cepheids with different period ranges: crosses are for Cepheids with periods shorter than 8 days (short-period Cepheid); triangles are for Cepheids with period in between 8 to 14 days (bump Cepheid); and filled circles are for Cepheids with periods longer than 14 days (long-period Cepheid). Error bars for each data point are omitted in the figure for clarity. The top two panels in Figure \\ref{fig:fig20} are Fourier intrarelations in V and I band respectively, and the bottom panel is the Fourier interrelations of $A_1(V)$ and $A_1(I)$.  It is clear from the figure that the scatter of $A_1-A_2$ intrarelations are larger than the scatter of $A_1(V)-A_1(I)$ interrelations. Furthermore, the short-period, bump and long-period Cepheids populate different regions in the plots of intrarelations, as is also seen in \\citet{ant87}. In contrast, The tightness of correlations in the Fourier interrelations is clear, and less dependent on period distribution. These two properties make Fourier interrelations more applicable in reconstructing the light curves than Fourier intrarelations.  \\subsection{Reconstructing Light Curves for Short Period Cepheids}  The relatively large ranges of the Fourier amplitudes for Cepheids with periods less than 10 days (or $\\log(P)<1.0$) is one reason for using the Fourier techniques to reconstruct the Cepheid light curves. From Figure \\ref{fig:fig3} to \\ref{fig:fig5}, it can be seen that the distribution of the Fourier amplitudes at periods longer than 10 days (or $\\log(P) > 1.0$) show certain trends, which can be used to construct templates light curves as a function of period. However, the Fourier amplitudes for Cepheids with period less than 10 days do not show any obvious trends but scatter around certain ranges (as given in Table \\ref{tab1}). For example, at a given short period, $A_1(V)$ may occupy the range from $\\sim 0.1$ to $\\sim 0.4$. Therefore, extra care has to be taken when applying template fitting to the Cepheids with periods less than 10 days. In fact, at periods shorter than $\\log(P) < 0.85$ the template techniques do not apply \\citep{ste96} but the Fourier techniques still holds good.  Although the extra-galactic Cepheids discovered in the past HST observations generally have period longer than 10 days, some short period period Cepheids have been discovered in Local Group galaxies IC 1613 \\citep{dol01} \\& Leo A \\citep{dol02} with HST and ground based observations, respectively. Hence for observing Cepheids at short period, the Fourier techniques are still useful to reconstruct the light curves. With the new installation of ACS (Advanced Camera for Survey) in HST, the discovery of more short period Cepheids in other galaxies is probable.  \\subsection{The Effect of Metallicity}  One motivation for studying direct Fourier techniques to reconstruct Cepheid light curves is to deal with the possible metallicity dependence in Cepheids. \\citet{ant00} shown that the metallicity may affect the shapes of light curves close to 10 days, based on the comparison of Cepheid light curves in two galaxies. The appropriate ranges of the Fourier parameters (as given in Table \\ref{tab1}) are assumed to cover the possible ranges due to the different metallicity environments. In addition, the Fourier interrelations are shown to be only weakly dependent on the metallicity though this does not mean that the actual light curve shape is independent of metallicity.  The template light curves defined by \\citet{ste96} are obtained from the Galaxy ($Z=0.020$), LMC ($Z=0.008$) and SMC ($Z=0.004$), with sample of $\\sim 30-45$ Cepheids in each galaxies. \\citet{ste96} looked for and failed to find differences between the Galactic, LMC and SMC Cepheids in terms of their amplitude and light curve shape sufficient to change the estimation of mean magnitudes. However, much more data is available now. \\citet{pac00} found statistically significant different amplitudes between the Galactic/LMC/SMC Cepheids in the period range $1.1 < \\log (P) < 1.4$, in the sense that higher metallicity Cepheids have higher amplitudes. Moreover, in the period range $0.0 < \\log (P) < 0.95$, the amplitudes ratio of $R_{21}$ and $R_{31}$ for Cepheids in LMC \\citep{and87} and SMC \\citep{and88} are found to be larger than the Galactic Cepheids. This may due to the different metallicity environments, at least for the case of the SMC \\citep{buc94}. Further some established workers in the field have found varying PC relations between the Galaxy, LMC and SMC and subsequently different PL relations \\citep{tam01,tam02}. Thus while it may be that template methods yield a sufficiently accurate mean in varying metallicity environments, we contend that the issue of light curve shape and metallicity needs to be revisited in the light of the data now currently available.  \\subsection{Applications in Determining the Cepheid Distances}  An immediate application of these Fourier techniques is to reconstruct the Cepheid light curves in NGC 4258 and derive its distance modulus. NGC 4258 is a spiral galaxy that has an accurate geometrical distance of $7.2\\pm0.3Mpc$ (corresponding to distance modulus of $29.28\\pm0.09$) from water maser measurement \\citep{her99}. By using the published data of 15 Cepheids discovered from HST observations \\citep{new01} and following the H0KP procedures (including using the same PL relations and reddening correction), we derived a distance modulus of $29.38\\pm0.06$ (random errors only) to NGC 4258, which is about $1.1\\sigma$ away from water maser distance. If we apply the $-0.07mag$ correction due to the WFPC2 calibration of ``long vs short'' exposures (this is caused by the charge-transfer efficiency of WFPC2, see H0KP and reference therein for details), then the final distance modulus becomes $29.31\\pm0.06$, which is consistent with the water maser distance. In addition, \\citet{kan02} have also derived the Cepheid distance to NGC 4258 with the same techniques presented here but slightly different PL relations, the result of $29.36\\pm0.06$ (random errors only) is also consistent to the water maser distance and the distance derived in this paper.  When we apply these techniques to determine Cepheid distances to 16 galaxies using published photometry by the H0KP, our distance moduli, on average, are very similar to H0KP results, with a difference of $\\sim0.01mag$ \\citep{kan02}. On the other hands, application of these techniques to 3 Sandage-Tammann-Saha galaxies (NGC 3627: \\citet{sah99}; NGC 3982: \\citet{sah01a}; NGC 4527: \\citet{sah01b}) yields shorter distance moduli of $0.07mags$ on average \\citep{kan02}. The good agreements between the results in \\citet{kan02} and in both H0KP and Sandage-Tammann-Saha galaxies suggest that our method is not too far of base."
    },
    "0212/astro-ph0212330_arXiv.txt": {
        "abstract": "We report initial results from the detection of optical emission in the arcsecond-scale radio jets of two quasars utilizing images from the {\\it Hubble Space Telescope} archive. The optical emission has a very knotty appearance and is consistent with synchrotron emission from highly relativistic electrons in the jet. Combining these observations with those of previously reported features in other quasars, an emerging trend appears to be that their radio-to-optical spectral indices are steeper than those of similar features in jets of lower power radio sources. ",
        "introduction": "\\label{}  The number of active galaxies observed to have radio-bright kpc-scale relativistic jets emanating from their nuclei has grown to many 100's over the last several decades \\citep{bri84,liu02}, but only a handful of them were known to have X-ray counterparts prior to the launch of the {\\it Chandra} X-ray Observatory.  {\\it Chandra} has in just the last three years netted X-ray detections of jet features in over 30 more radio sources \\citep[see][and the most recent updates in the associated website -- http://hea-www.harvard.edu/XJET/]{harr02}, and many more are being discovered as a result of dedicated searches (see e.g.  individual contributions by R.~M. Sambruna, H.~L. Marshall, and D.~A. Schwartz). With the increased interest in the X-rays, studies in the optical/UV regime have also enjoyed a resurgence of activity. This is because the optical is the only other accessible window in the electromagnetic spectrum where comparable sub-arcsecond resolution can be readily achieved with the {\\it Hubble Space Telescope}\\footnote{Based on observations made with the NASA/ESA Hubble Space Telescope, obtained from the data archive at the Space Telescope Science Institute. STScI is operated by the Association of Universities for Research in Astronomy, Inc. under NASA contract NAS 5-26555.}, and bridges the $\\sim$6-7 orders of magnitude gap in frequency coverage between radio and X-ray studies of resolved features. It turns out that the optical fluxes and spectra can be a key diagnostic in distinguishing between different X-ray production mechanisms.  Specifically, the faintness of the optical knots in the `first-light' X-ray jet detection in the powerful quasar PKS~0637--752 \\citep{cha00,sch00} helped to rule out canonical synchrotron and synchrotron self-Compton models, and established the importance of X-rays produced via inverse Compton scattering of CMB photons by synchrotron emitting electrons in the kpc-scale jets \\citep{tav00,cel01}. Also, if the emission mechanism responsible for the optical emission is synchrotron radiation, as is commonly accepted, observations in this waveband probe the physics of the most energetic and short-lived electrons, and are key to studying acceleration processes in the jet.  In the pre-{\\it Chandra} era, optical jets were predominantly found in low power sources with the {\\it HST} \\citep[e.g.,][]{spa94,sca02}, and optical counterparts to jets in quasars were difficult to detect with ground based facilities. Only a handful of examples existed, courtesy of the {\\it HST} \\citep[e.g.,][]{rid97}, so optical jets in quasars were thought to be rare. However, {\\it Chandra's} successful X-ray detections of quasar jets has forced observers to examine new and archival {\\it HST} data more carefully, resulting in many new detections in the optical.  This resulted in the identifications of a number of bright optical knots coincident with peaks in the radio jets of about ten quasars in a recent {\\it Chandra/HST} cycle-2 search for such emission \\citep[][R. Scarpa et al., in prep.]{sam02}. We applied the same techniques employed in this survey and found a bright optical/UV knot in the kpc-scale radio jet of 3C~279 in archival {\\it HST} data \\citep{che02} which encouraged us to search for more such emission in other available archival data. ",
        "conclusions": ""
    },
    "0212/astro-ph0212025_arXiv.txt": {
        "abstract": "Heise et al. (2001) have reported the identification of a new class of sources by Beppo-SAX, the so-called X-ray flashes (XRFs), which have many common properties with $\\gamma$-ray bursts (GRBs) but are not detected in the $\\gamma$-ray range. In the framework of the internal shock model, we investigate the possibility that XRFs have the same origin than GRBs but are intrinsically softer due to different shock parameters. ",
        "introduction": "\\vspace*{-1ex}  Heise et al. (2001) define XRFs with 4 criteria : (i) a strong non-thermal X-ray emission; (ii) a weak $\\gamma$-ray emission; (iii) a duration less than a few 1000 s; (iv) no strong quiescent V/IR counterpart. In 5 years, Beppo-SAX detected 17 XRFs. Their distribution over the sky seems to be isotropic, their durations are comparable with those of GRBs detected by Beppo-SAX but their X- to $\\gamma$-ray fluence ratios are higher. Since its launch, HETE-2 has also detected a few  XRFs. The preliminary study confirms that these events have many common properties with GRBs, except the fact that they are not detected above 50--100 keV.\\\\ \\vspace*{-3.75ex} ",
        "conclusions": "\\vspace*{-1ex}  Our main results are : (i) a population of XRFs is indeed obtained if one adopt a broad range for the IS parameters : $\\dot{E}$, $t_\\mathrm{var}$, $K$, $\\Gamma_\\mathrm{min}$ and $C$; (ii) among all the possibilities listed above, \\textit{small contrasts of the Lorentz factor} (low $C$) seem to produce the better agreement with the observed properties of XRFs. The conclusion is that \\textit{in the framework of the IS model, XRFs are most likely soft GRBs associated with smooth outflows}. A possible observational test is that XRFs should have ``normal'' afterglows, but a low ratio of the afterglow over prompt energy outputs. These results will be presented in details in a forthcoming paper.\\\\ \\vspace*{-3.75ex}"
    },
    "0212/astro-ph0212213_arXiv.txt": {
        "abstract": "We present the results of two extensive {\\it Rossi X-ray Timing Explorer} observations of large X-ray flaring episodes from the high-mass X-ray binary pulsar LMC X-4. Light curves during the flaring episodes comprise bright peaks embedded in relatively fainter regions, with complex patterns of recurrence and clustering of flares. We identify precursors preceding the flaring activity. Pulse profiles during the flares appear to be simple sinusoids, and pulsed fractions are proportional to the flare intensities. We fit Gaussian functions to flare peaks to estimate the mean full-width-half-maximum to be $\\sim$68 s. Significant rapid aperiodic variability exists up to a few hertz during the flares, which is related to the appearance of narrow, spiky peaks in the light curves. While spectral fits and softness ratios show overall spectral softening as the flare intensity increases, the narrow, spiky peaks do not follow this trend. The mean fluence of the flare peaks is (3.1 $\\pm$ 2.9) $\\times$ 10$^{40}$ ergs in the 2.5--25 keV energy range, with its maximum at $\\sim$1.9 $\\times$ 10$^{41}$ ergs. The flare peak luminosity reaches up to (2.1 $\\pm$ 0.2) $\\times$ 10$^{39}$ ergs s$^{-1}$, far above the Eddington luminosity of a neutron star. We discuss possible origins of the flares, and we also propose that inhomogeneous accretion columns onto the neutron star polar caps are responsible for the observed properties. ",
        "introduction": "Out of $\\sim$200 known neutron star/black hole X-ray binaries, more than 40 sources have been observed with a wide range of coherent spin period of $\\sim$0.07--1400 s, identifying the central source as a rotating, highly magnetized neutron star powered by accretion (van Paradijs 1995; Bildsten et al. 1997). These accretion-powered X-ray binary pulsars (hereafter, X-ray pulsars) often show abrupt increases in X-ray luminosity known as ``bursts\" or ``flares.\" About 30 \\% of the X-ray pulsars have been observed to show bursts and/or flares, including (1) the low-mass X-ray binary 4U 1626--67 (McClintock et al. 1980), (2) the bursting pulsar GRO J1744--28 (Fishman et al. 1995; Kouveliotou et al. 1996), (3) the high-mass X-ray binaries such as LMC X-4 (e.g., Levine et al. 1991; Woo, Clark, \\& Levine 1995; Levine, Rappaport, \\& Zojcheski 2000; Moon \\& Eikenberry 2001) and SMC X-1 (Angelini, Stella, \\& White 1991; Moon, Eikenberry, \\& Wasserman 2002), and (4) the transient Be binary system EXO 2030+375 (Parmar et al. 1989). Although they are easily identified as sudden increases of X-ray luminosity in most cases, the detailed properties of the bursts/flares differ much from source to source. For instance, whereas bursts from the high-mass system SMC X-1 have been detected only once despite extensive studies of the source \\footnote{SMC X-1 has been identified as a flaring X-ray pulsar very recently (Moon et al. 2002).}, another high-mass system LMC X-4 has been observed with large, semi-regular flaring episodes over the last two decades by various X-ray satellites. This diversity in bursts/flares suggests that many different mechanisms are responsible for the observed phenomena.  Most theoretical studies of the origin of bursts/flares from X-ray pulsars have relied upon one of three following mechanisms: (1) accretion disk instabilities, (2) nuclear burning on the surface of the neutron star, and (3) non-uniform stellar winds from mass-donating companions. Obviously the most well known case for the accretion disk instability is the X-ray bursts from the bursting pulsar GRO 1744--28, for which a thermal-viscous instability in the accretion disk has been investigated to explain the origin of the bursts (e.g., Cannizzo 1996, 1997). The X-ray bursts from GRO 1744--28 share many properties with the Type II X-ray bursts from the Rapid Burster (Lewin et al. 1996) and the X-ray flares from SMC X-1 (Moon et al. 2002). On the other hand, the presence of a He-rich dwarf companion led Brown \\& Bildsten (1998) to suggest the possibility that $\\sim$1000-s flares from 4U 1626--67 are caused by carbon burning under the neutron star surface. Taam, Fryxell, \\& Brown (1988) attributed the origin of the flares from EXO 2030+375 to an accretion disk instability caused by the inhomogeneity in the high-velocity ($\\sim$550 km s$^{-1}$) winds from a Be star companion.  LMC X-4 is a persistent, disk-fed, high-mass system with a pulsational and an orbital period of $\\sim$13.5 s and $\\sim$1.4 d, respectively (Kelley et al. 1983). It is one of a few X-ray binaries showing a super-orbital period ($\\sim$30.3 day; Ilovaisky et al. 1984) possibly caused by a precessing, tilted accretion disk. The optical companion is a 14th magnitude O-type star of $\\sim$15 $M_{\\odot}$ (Sanduleak \\& Philip 1977). Over the last two decades, LMC X-4 has exhibited large X-ray flaring episodes which have made it one of the most regular and representative flaring X-ray sources. During the flares, the X-ray spectrum softens considerably, and the pulse profiles become simple, sinusoidal (Levine et al. 1991, 2000). Recently, Moon \\& Eikenberry (2001) reported two relatively long-time scale structures in the flares: quasi-periodic oscillations of $\\sim$0.65--1.35 and $\\sim$2--20 mHz. They also found that the amplitudes of the flares are almost exactly modulated by those of the neutron star's coherent pulsations (or vice versa), indicating that the flares occur near the magnetic pole of the neutron star. Although the LMC X-4 flares have been observed for many years with various X-ray satellites, no convincing mechanism has been proposed to explain the origin of the flares (e.g., Levine et al. 2000).  On the other hand, LMC X-4 flares have been known as super-Eddington phenomena: sometimes the flare X-ray luminosity reaches up to $\\sim$10$^{39}$ ergs s$^{-1}$ (while its normal state luminosity is $\\sim$2 $\\times$ 10$^{38}$ ergs s$^{-1}$, comparable to the Eddington luminosity of the $\\sim$1.4 $M_{\\odot}$ neutron star of LMC X-4). Because the key physical parameters of LMC X-4 -- such as mass, distance, and magnetic field  -- are relatively well determined to be 1.4 $\\pm$ 0.3 $M_{\\odot}$ (Levine et al. 1991), 50 $\\pm$ 10 kpc (e.g., Kov\\'{a}cs 2000), and $\\sim$10$^{13}$ G (La Barbera et al. 2001), respectively, the LMC X-4 flares may provide a rare example for studying super-Eddington phenomena thoroughly without assuming the key parameters. This is not the case for most super-Eddington sources owing to relatively large uncertainties in their distances and masses. Considering the recent important controversy over the existence of the so-called ``intermediate-mass black holes'' (e.g., King et al. 2001; Begelman 2001, 2002), for which mass ($\\sim$10$^2$--10$^4$ $M_{\\rm \\odot}$) is determined under the assumption that the observed X-ray luminosity is the Eddington luminosity, the investigation of the origin of the super-Eddington radiation from the LMC X-4 flares can also be useful to investigate the origin of the claimed intermediate-mass black holes.  In this paper, we present by far the most extensive observations of the LMC X-4 flares made with the {\\it Rossi X-ray Timing Explorer} (RXTE). We describe the RXTE observations and our data analyses in \\S 2; we present the results of timing and spectral analyses in \\S 3. We discuss the possible origin of the LMC X-4 flares, as well as their super-Eddington radiation, in \\S 4, We summarize our results in \\S 5. ",
        "conclusions": "In this paper, we have analyzed two extensive observations of the LMC X-4 flares made with RXTE. We summarize our main results as follows. (1) The flare light curves reveal the existence of the flat precursors preceding the flares, with the complex activity of the LMC X-4 flares. (2) The FWHM and mean fluence of the flares are 68 $\\pm$ 31 s and (3.1 $\\pm$ 2.9) $\\times$ 10$^{40}$ ergs, respectively. (3) The pulse profiles during the flares are broad, sinusoidal, and the pulsed fractions increase with the flare intensities. (4) With 0.0625-s resolution, narrow, spiky peaks appear in the light curves, so does the rapid aperiodic variability in PDSs. Both intensify along with the flare intensities. (6) The softness ratios and spectral fits show that the flare spectra soften as the intensities increase. (7) With 0.0625-s resolution, while the photons having count rates less than 300 counts/sec/PCU show the softness ratio distributions almost equal to those obtained with 32-s resolution, the photons having larger photon count rates show a much slower increase in the softness ratios as the flare intensities increase. (8) The peak luminosity of the flares reaches up to $\\sim$2.1 $\\times$ 10$^{39}$ ergs s$^{-1}$, well above the Eddington limit for an 1.4 $M_{\\odot}$ neutron star.  We find it difficult to explain the observed properties of the LMC X-4 flares in terms of the thermonuclear instability on the neutron star surface or the viscous instability at the inner accretion disk. We consider the possibility that the plasma Rayleigh-Taylor instability in the accretion disk may be responsible for the flares. We propose that increased optical depths and density inhomogeneity in the accretion columns onto the neutron star polar cap during the LMC X-4 flares are responsible for the observed properties, including the super-Eddington luminosities of the flares. If confirmed, this is the first observational evidence for the inhomogeneous atmosphere responsible for super-Eddington radiation. The broad pulse profiles of the LMC X-4 flares may be caused by substantial gravitational bending of the flare emission transported via fan beams."
    },
    "0212/astro-ph0212163_arXiv.txt": {
        "abstract": "Previous investigations on hydrogen-rich white dwarfs generally yield only very small rotational velocities ($v_{\\mathrm{rot}}\\cdot\\sin i$). We have analyzed line profiles in high-resolution optical spectra of eight hydrogen-deficient (pre-) white dwarfs and find deviations from the dominant Stark line broadening in five cases which, interpreted as an effect of stellar rotation, indicate projected rotational velocities of $40 - 70 \\mathrm{km/sec}$. For the three least luminous stars upper limits of $v_{\\mathrm{rot}}\\cdot\\sin i = 15 - 25 \\mathrm{km/sec}$ could be derived only. The resulting velocities correlate with luminosity and mass. However, since the mass-loss rate is correlated to the luminosity of a star, the observed line profiles may be affected by a stellar wind as well. In the case of RX\\,J2117.1+3412\\index{RX J2117.1+3412}, this would solve discrepancies to results of pulsational modeling ($v_{\\mathrm{rot}}\\cdot\\sin i \\approx 0$). ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/astro-ph0212480_arXiv.txt": {
        "abstract": " ",
        "introduction": "It is now generally accepted that the subluminous B (sdB) stars can be identified with models for Extreme Horizontal Branch (EHB) stars burning He in their core, but with a very tiny inert hydrogen envelope (Heber, 1986). An EHB star bears great resemblance to a helium main-sequence star of half a solar mass and its further evolution should proceed similarly (i.e. directly to the white dwarf graveyard) as confirmed by evolutionary calculations (Dorman et al. 1993). How stars evolve to the EHB configuration is controversial. The problem is how the mass loss mechanism in the progenitor manages to remove all but a tiny fraction of the hydrogen envelope at {\\em precisely} the same time as the He core has attained the minimum mass ($\\approx0.5$M$_\\odot$) required for the He flash.  There is growing evidence that close binary evolution is an important if not the dominant formation path of sdB stars (Maxted et al. 2001a, Heber et al. 2002). SdB stars can result from stable Roche lobe overflow or common envelope ejection models (Han et al. 2002). The companions are either low mass main sequence stars or white dwarfs. SdB stars may also result from the merger of two He white dwarfs (Webbink, 1984). ",
        "conclusions": ""
    },
    "0212/astro-ph0212449_arXiv.txt": {
        "abstract": "We point out that two problems of observational cosmology, the facts i) that $\\ga 60 \\%$ of the baryonic content of the universe is not observed at $z\\sim 0$ and ii) that the properties of small clusters do not agree with simple expectations, could be closely related. As shown by recent studies, the shock-heating associated with the formation of large-scale structures heats the intergalactic medium (IGM) and leads to a ``warm IGM'' component for the gas. In the same spirit, we suggest the intracluster medium (ICM) to be a mixture of galaxy-recycled, metal enriched gas and intergalactic gas, shock-heated by the collapsing much larger scales. This could be obtained through two processes: 1) the late infalling gas from the external warm IGM is efficiently mixed within the halo and brings some additional entropy, or 2) the shocks generated by larger non-linear scales are also present within clusters and can heat the ICM. We show that if assumption (1) holds, the entropy brought by the warm IGM is sufficient to explain the observed properties of clusters, in particular the entropy floor and the $\\LX-T$ relation. On the other hand, we briefly note that the scenario (2) would require a stronger shock-heating because of the larger density of the ICM as compared with filaments. Although the efficiency of these two processes remains to be checked on a quantitative level, they present the advantage to dispense with the need to invoke any strong preheating from supernovae or quasars (which has otherwise been introduced for the sole purpose of reproducing the behaviour of clusters). Matter ejection by galaxies is included in the present calculations, and consistently with the metal-enrichment requirements, is indeed shown to yield only a quite moderate entropy increase. Our scenario of clusters being \"born warm\" can be checked through the predicted redshift evolution of the entropy floor. ",
        "introduction": "A key test of cosmological scenarios is to check whether they can reproduce the large-scale structure of the present universe. Many studies have already shown that the usual hierarchical scenarios (such as  the standard CDM model) agree reasonably well with observations of galaxies (Valageas \\& Schaeffer 1999, Kauffmann et al. 1993, Cole et al. 1994), Lyman-$\\alpha$ clouds (Valageas et al. 1999, Petitjean et al. 1992, Miralda-Escude et al. 1996), and quasars, as well as with  reionization constraints (Valageas \\& Silk 1999a, Gnedin \\& Ostriker 1997, Haiman \\& Loeb 1998).  However some puzzles still remain unresolved. For instance, the mass of baryons observed at $z=0$ (stars, intracluster hot gas) yields $\\Ob \\sim 0.01$ (Fukugita et al. 1998), while the Lyman alpha forest accounts for about as much (Penton et al. 2000). This falls short of the CMB determination by Archeops (Benoit et al. 2003), $\\Ob \\simeq 0.052 \\pm 0.007$, or the prediction of standard nucleosynthesis $\\Ob \\simeq 0.045 \\pm 0.006$ (Tytler et al. 2000) \\footnote{Our choice of cosmological parameters is given in Section \\ref{Phase-diagram of cosmological baryons}.}. As argued in Cen \\& Ostriker (1999) and Dave et al. (2001), up to $50\\%$ of the baryons could be in the form of a ``warm\" phase of the intergalactic medium (IGM), with temperatures in the range $10^5 < T <10^7$ K, and provide the explanation for the ``missing\" baryonic matter.  A second problem is the temperature - X-ray luminosity relation of clusters. Indeed, the simplest models yield $\\LX \\propto T^2$ (Kaiser 1986) which disagrees with observations (Ponman et al. 1996). It has been suggested (e.g., Evrard \\& Henry 1991, Cavaliere et al. 1997, Ponman et al. 1999) that this discrepancy could be due to a ``preheating'' of the IGM which would raise its entropy before clusters form. Such an entropy ``floor'' breaks the standard scaling laws and  yields a steeper  $T - \\LX$ relation, and may possibly  play a  role in the solution of the ``overcooling'' problem (Blanchard et al. 1992). Two sources of energy which could provide this preheating appear naturally in the usual cosmological scenarios: supernovae/stellar winds and quasars. However, realistic models (e.g., Valageas \\& Silk 1999b, Wu et al. 2000, Bower et al. 2001) show that the energy released by supernovae is unlikely to be sufficient (unless the energy transfer is highly efficient and targeted to the gas which will build clusters and groups). On the other hand, the energy provided by quasars is largely  sufficient (Valageas \\& Silk 1999b), but it is not easy to estimate with a good accuracy the efficiency of the energy transfer. Although it is commonly agreed that an increase in entropy is required, its actual origin remains an open issue.  Some numerical simulations (Muanwong et al. 2001), on the other hand,  show that inhomogeneous cooling of the baryonic gas, which removes the cooler gas as it condenses into galaxies while leaving only the hotter component available as intracluster gas, can play the same role as an increase in entropy. A related proposal (Voit \\& Bryan 2001) is that the entropy along the edge of the cooling curve matches, at a given virial temperature, the observed entropy of the intracluster gas. However, to be efficient this may imply an excessive amount of cooling (Muanwong et al. 2001).  In this work, we investigate another source of entropy: the shock-heating associated with the gravitational collapse of very large scales which are just turning non-linear. This is partly motivated by the recent finding in numerical simulations of a warm IGM phase generated by such a {\\it transfer of energy from collapsing very large scales to much lower scales}, see Dave et al. (2001). This mechanism, driven by the non-linear length scale ($\\sim 9$ Mpc at $z=0$), heats the IGM which is clustered within much smaller filaments (i.e. with a similar length but a smaller thickness of order of a few hundred kpc at $z=0$) and occupies a small fraction of the volume of the universe (e.g., Cen \\& Ostriker 1999, Valageas et al. 2002). Since clusters are located at the crossing of such filaments, their properties could also be influenced by this shock-heating. Therefore, in this study we investigate whether the entropy floor observed within clusters could be generated by this shock-heating from scales much above the cluster scale. This could occur through two mechanisms. First, the warm IGM embedded within neighbouring filaments, where it underwent this shock-heating, could be efficiently mixed within the cluster as the accretion which builds the halo proceeds. Then, the additional entropy brought with this flow of matter could increase the entropy of the cluster. In order to be relevant, this process requires a sufficiently efficient mixing of the gas so that at least the small clusters with a temperature below $\\sim 1$ keV which display the anomaly are affected in their bulk gas distribution. We shall evaluate the level of mixing required by such a scenario. A second mechanism would be for the gas to be shock-heated {\\it in situ}, that is the shocks generated by the large scales which are just turning non-linear may run through the cluster and heat the ICM itself. We shall not study this alternative scenario in detail here, but we note that it requires stronger shocks in order to significantly increase the entropy of this higher-density gas. ",
        "conclusions": "\\label{Conclusion}  We have shown that the entropy of the warm IGM can be sufficient to explain the observed entropy of small clusters. This scenario relies on the assumption that the gas which has recently falled into the cluster from the surrounding filaments has been efficiently mixed within the halo so as to form at least $28 \\%$ of the ICM at the radii where the extra-entropy is observed. This point is clearly a matter of debate since it is not obvious to follow the behaviour of the gas within such complex objects as clusters. However, we feel this requirement is quite plausible. A natural scale appears in our scenario, $1$ keV, the potential energy of the scales just turning non-linear. Clusters below $1$ keV where this external entropy source can play a role are specifically the ones under debate. These rather small clusters should be rather irregular so that their accretion is not expected to be well-ordered or spherically symmetric. Then, the late gas coming from the warm IGM may indeed spread over the whole group, taking advantage of the general mixing (e.g., turbulence) brought by the chaotic dynamics of the various components (e.g., galaxies) of the group. On the other hand, we briefly notice that one could also imagine a second scenario. There, the gravitational energy associated with the large scales which are just turning non-linear could directly heat the ICM itself. This assumes that shocks generated by those large scales can spread within the cluster and heat at least $28 \\%$ of the ICM. However, in order to reach the same entropies this scenario requires a stronger shock-heating than in the first model because of the need to compensate for the higher density of the gas within the ICM as compared with filaments.  Altough the efficiency of these mechanisms remains to be evaluated with a better accuracy before one can safely determine whether they indeed play an important role in the physics of clusters, we note that such scenarios present the advantage to dispense with the need of a new external preheating source (e.g., supernovae or quasars). Indeed, they only rely on gravitational shock-heating which has recently been seen to explain the properties of the IGM in numerical simulations (e.g., Dave et al. 2001). Therefore, we merely propose here that the mechanisms which build the warm IGM may also govern the properties of small clusters. So far, numerical simulations have not really substantiated this claim since they usually recover the standard scaling $L_X \\propto T^2$, unless one adds some ad-hoc preheating or takes into account cooling processes (which lead to other problems like overcooling). However, they may not have been able to pin down all salient effects of shock-heating mechanisms yet.  Similar processes may be expected to be at work at the epoch of galaxy formation. They may for instance modify the hydrodynamical equilibrium of the smaller galaxies. In particular, as already noticed in Valageas \\& Silk (1999b) preheating models are strongly constrained by the contradictory requirements arising from the observed properties of clusters and galaxies. This will be studied in a forthcoming paper for the process presented here which exhibits a specific evolution with redshift.  In summary, there are three hypothesis for the entropy floor in a given cluster. It may arise 1) from external preheating by supernovae or quasars; 2) from a very specific feature of the cooling curve, as implied by the suggestion of Voit \\& Bryan (2001); 3) as we suggest here, from the pre-collapse phase of structures much larger than clusters which transfer their energy to the smaller cluster scales through shock-heating, which also builds the warm IGM phase.  We have stressed the difficulties encountered by the two first possibilities. The advantage of the latter explanation is that it is a natural outcome of hierarchical scenarios of large-scale structure formation. Numerical simulations and observations may be used to check the entropy evolution implied by our suggestion."
    },
    "0212/astro-ph0212355_arXiv.txt": {
        "abstract": "Hydrodynamic cosmological simulations predict that the average opacity of the Ly$\\alpha$ forest should increase in the neighborhood of galaxies because galaxies form in dense environments.  Recent observations (Adelberger et al. \\cite{adelberger02}) confirm this expectation at large scales, but they show a {\\it decrease} of absorption at comoving separations $\\Delta_r \\ltorder 1 h^{-1}$ Mpc.  We show that this discrepancy is statistically significant, especially for the innermost data point at $\\Delta_r \\le 0.5 h^{-1}$ Mpc, even though this data point rests on three galaxy-quasar pairs.  Galaxy redshift errors of the expected magnitude are insufficient to resolve the conflict.  Peculiar velocities allow gas at comoving distances $\\gtorder 1 h^{-1}$ Mpc to produce saturated absorption at the galaxy redshift, putting stringent requirements on any ``feedback'' solution.  Local photoionization is insufficient, even if we allow for recurrent AGN activity that keeps the neutral hydrogen fraction below its equilibrium value.  A simple ``wind'' model that eliminates all neutral hydrogen in spheres around the observed galaxies can marginally explain the data, but only if the winds extend to comoving radii $\\sim 1.5 h^{-1}$ Mpc. ",
        "introduction": "In a recent paper \\cite{kollmeier02}, we discuss a variety of predictions for galaxy-Ly$\\alpha$ forest correlations from hydrodynamic simulations. In this proceeding, we focus on the ``galaxy proximity'' effect on small scales ($\\le 2 h^{-1}$ Mpc comoving).  Using smoothed particle hydrodynamics simulations (SPH) of a $\\Lambda$CDM universe ($\\Omega_m=0.4, \\Omega_\\Lambda=0.6, h=0.65,\\Omega_b=0.02 h^{-2}=0.0473$, $\\sigma_8=0.80$), we generate synthetic Ly$\\alpha$ forest spectra from the temperature, gas density, and velocity at each spatial location in skewers through the simulation box.  Once we have created synthetic spectra, we use the known positions of the simulated galaxies to compute the mean flux decrement, $\\langle D\\rangle = \\langle 1 - e^{-\\tau}\\rangle$, as a function of comoving separation from the galaxy, $\\Delta_r$.  Figure 1a shows the basic predictions for the mean flux decrement, computed for the 150 galaxies with the highest star formation rates within a simulation $50 h^{-1}$ Mpc on a side.  We see a clear trend of increasing decrement (absorption) with decreasing $\\Delta_r$, a signature of the dense environments these galaxies occupy. The observed points, from \\cite{adelberger02}, show a similar trend at large scales, but the trend flattens at $\\Delta_r \\ltorder 2.5 h^{-1}$ Mpc, and absorption {\\it decreases} for $\\Delta_r \\ltorder 1 h^{-1}$ Mpc.  The innermost point, $0 \\le \\Delta_r \\le 0.5 h^{-1}$ Mpc is in especially severe conflict with the theoretical prediction.  We investigate several possibilities for resolving this conflict below.  Similar investigations have been carried out by \\cite{croft02}, \\cite{bruscoli02}, and our results are compatible with theirs to the extent that they overlap.  \\begin{figure} \\includegraphics[height=.5\\textheight]{fig1.eps} \\caption{The Conditional Mean Flux Decrement.  (a) The conditional mean decrement compared with the baseline calculation, and calculations with redshift errors as indicated in panel and described in text, (b) Effect of photoionization feedback from stars or AGN on the numerical predictions, (c) Effect of spherical winds/gas removal on the predictions, (d) Cumulative distribution of 3-tuples at $\\Delta_r=0.25$ Mpc as a function of $\\langle D\\rangle$. Vertical lines at $\\langle D\\rangle=0.11$ and $0.36$ correspond to the Adelberger et al. \\cite{adelberger02} value for this separation and the global mean respectively. Symbols are as indicated in the panel.} \\end{figure} ",
        "conclusions": ""
    },
    "0212/astro-ph0212119_arXiv.txt": {
        "abstract": "\\noindent For the first time, {\\it all} available pseudo-Schwarzschild potentials are exhaustively used to investigate the possibility of shock formation in hydrodynamic, invicid, black hole accretion discs. It is shown that a significant region of parameter space spanned by important accretion parameters allows shock formation for flow in {\\it  all} potentials used in this work.  This leads to the conclusion that the standing shocks are essential ingredients in accretion discs around non-rotating black holes in general. Using a complete general relativistic framework, equations governing multi-transonic black hole accretion and wind are also formulated and solved in the Schwarzschild metric. Shock solutions for accretion flow in various pseudo potentials are then compared with such general relativistic solutions to identify which potential is the best approximation of Schwarzschild space-time as far as the question of shock formation in black hole accretion discs is concerned. ",
        "introduction": "\\noindent The process by which any gravitating, massive, astrophysical object captures its surrounding fluid is called accretion. Depending on the rotational energy content of the infalling material, accretion flows onto  black holes may be broadly classified into two different categories, i.e., non-rotating (spherical)  and rotating accretion (accretion discs). If the instantaneous dynamical velocity and local acoustic velocity of the accreting fluid, moving along a space curve parameterized by $r$, are $u(r)$ and $a(r)$ respectively, then the local Mach number $M(r)$ of the fluid can be defined as $M(r)=\\frac{u(r)}{a(r)}$. The flow will be locally subsonic or supersonic according to $M(r) < 1$ or $ >1$, i.e., according to $u(r)<a(r)$ or $u(r)>a(r)$. The flow is transonic if at any moment it crosses $M=1$. This happens when a subsonic to supersonic or supersonic to subsonic transition takes place either continuously or discontinuously. The point(s) where such crossing takes place continiously is (are) called sonic point(s), and where such crossing takes place discontinuously are called shocks or discontinuities. It is generally argued that, in order to satisfy the inner boundary conditions imposed by the event horizon, accretion onto black holes exhibit transonic properties in general, which further indicates that formation of shock waves are  possible in astrophysical fluid flows onto galactic and extra-galactic black holes. One also expects that shock formation in black hole accretion might be a general phenomena because shock waves in rotating and non-rotating flows are convincingly able to provide an important and efficient mechanism for conversion of significant amount of the gravitational energy (available from deep potential wells created by these massive compact accretors) into radiation by randomizing the directed infall motion of the accreting fluid. Hence shocks possibly play an important role in governing the overall dynamical and radiative processes taking place in accreting plasma. Thus the study of steady, stationary shock waves produced in black hole accretion has acquired a very important status in recent years and it is expected that shocks may be an important ingredient in an accreting black hole system in general. \\\\ \\noindent While the possibility of the formation of a standing spherical shock around compact objects was first conceived long ago (Bisnovatyi-Kogan, Zel`Dovich, \\& Sunyaev 1971), most of the works on shock formation in spherical accretion share more or less the same philosophy that one should incorporate shock formation to increase the efficiency of directed radial infall in order to explain the high luminosity of AGNs and QSOs and to model their broad band spectrum (Jones \\& Ellison 1991). Considerable work has been done in this direction where several authors have investigated the formation and dynamics of standing shock in spherical accretion (M\\'esz\\'aros \\& Ostriker 1983, Protheros \\& Kazanas 1983, Chang \\& Osttriker 1985, Kazanas \\& Ellision 1986, Babul, Ostriker \\& M\\'esz\\'aros 1989, Park 1990, 1990a). Ideas and formalisms developed in these works have been applied to study related interesting problems like entropic-acoustic or various other instabilities in spherical accretion (Foglizzo \\& Tagger 2000, Blondin \\& Ellison 2001, Lai \\& Goldreich 2000, Foglizzo 2001, Kovalenko \\& Eremin 1998), production of high energy cosmic rays from AGNs (Protheroe \\& Szabo  1992), study of the hadronic model of AGNs (Blondin  \\& Konigl 1987, Contopoulos \\& Kazanas 1995), high energetic emission from relativistic particles in our galactic centre (Markoff, Melia \\& Sarcevic 1999), explanation of high lithium abundances in the late-type, low-mass companions of the soft X-ray transient, (Guessoum \\& Kazanas 1999), study of accretion powered spherical winds emanating from galactic and extra galactic black hole environments (Das 2001).\\\\ \\noindent With equal (if not more) importance and rigor, the question of shock formation in accretion discs around Schwarzchild black holes has been addressed by several authors. While the initial works in this direction can be attributed to Fukue (1983),  Hawley, Wilson \\& Smarr (1984), Ferrari et al. (1985),  Swada et al.  (1986) and Spruit (1987), it was Fukue (1987) and Chakrabarti and his collaborators (Chakrabarti 1989 (C89 hereafter), 1996 and references therein, Abramowicz \\& Chakrabarti 1990, Chakrabarti \\& Molteni 1993) who were the first to provide the satisfactory semi-analytical or numerical global shock solution for transonic, invicid, Keplerian or sub-Keplerian rotating accretion around a Schwarzchild black hole. Consequently, their works were further supported and improved by several other independent works (Yang \\& Kafatos 1995 (YK hereafter), Caditz \\& Tsuruta 1998, T\\'oth, Keppens \\& Botchev 1998). Because of the inner boundary conditions imposed by the event horizon, shocks form in BH accretion discs only if the flow has more than one real physical $X$ type sonic point (multi-transonic flow). For a particular set of initial boundary conditions, some of the above mentioned works report multiplicity in shock location, but such a degeneracy can ultimately be removed by local stability analysis, allowing one to assert that only one stable shock location is possible. Hereafter, whenever we will use the word `shock', it is to be understood that we will, in general, always refer only the stable shock location unless otherwise mentioned.\\\\ \\noindent The above mentioned works deserve attention because the shocked flows studied there are expected to explain the spectral properties of BH candidates. However, thus far in the astrophysical literature, the theoretical study of steady, standing shock formation in accretion discs around non-rotating BHs has suffered from two general limitations. Firstly, the shock solutions were obtained either on a case by case basis, or, even when successful attempts were made  to provide a more complete analysis, the boundary of the parameter space responsible for shock formation was obtained only for global variation of the total specific energy ${\\cal E}$ (or accretion rate ${\\dot M}$) and specific angular momentum $\\lambda$ of the flow, and not for variations of the polytropic constant $\\gamma$ of the flow; rather accretion was always considered to be ultra-relativistic \\footnote{By the term `ultra-relativistic' and `purely non-relativistic' we mean a flow with $\\gamma=\\frac{4}{3}$ and $\\gamma=\\frac{5}{3}$ respectively, according to the terminology used in Frank et. al. 1992.} which may {\\it not} always be a realistic assumption. As $\\gamma$ is expected to have great influence on the radiative properties of the flow in general, we think that ignoring the explicit dependence of shock solutions on $\\gamma$ limits claims of generality. Secondly, except for YK, all available so called global shock solutions have been discussed only in the context of one particular type of BH potential, namely, the Paczy\\'nski \\& Wiita (1980) potential ($\\Phi_1$ hereafter). Along with the $\\Phi_1$, recent studies (Das \\& Sarkar 2001; hereafter DS, and references therein) enhance the importance of also considering three other pseudo-Schwarzschild BH potentials, one ($\\Phi_2$ hereafter) proposed by Nowak \\& Wagoner (1991), and two others ($\\Phi_3$ and $\\Phi_4$ hereafter) due to Artemova, Bj\\\"{o}rnsson \\& Novikov (Artemova, Bj\\\"{o}rnsson \\& Novikov 1996, ABN hereafter), in mimicking the complete general relativistic space-time for accretion around a Schwarzschild black hole. Hence we believe that being restricted to only one specific pseudo-Schwarzschild BH potential does  not guarantee the claimed `global' nature of so called global shock solutions present in the literature, rather one must study the transonic disc structure as well as shock formation in {\\it all} available BH potentials to firmly assert the ubiquity of shock formation in multi-transonic accretion disc around a Schwarzschild BH. In this context, it is to be mentioned here that YK deserves special importance because the shock solution due to YK appears to be the only work available in the literature which provides the complete general relativistic description of shock formation {\\it exclusively} for a non-rotating BH. Nevertheless, this work deals only with isothermal accretion but one understands that global isothermality in BH accretion is difficult to achieve for realistic flows and more general kind of BH accretion is expected to be governed by polytropic equation of state. Also YK does not provide the global parameter space dependence of shock solutions. A few authors claim to provide the full general relativistic shock solutions for Schwarzschild BHs as a limiting case of their results obtained in Kerr geometry (Chakrabarti 1996a,b; Lu, Yu, Yuan \\& Young 1997, LY3 hereafter). In doing so, a number of assumptions were made, some of which, however, may not appear to be fully convincing. For example, either the disc was supposed to be in conical equilibrium (LY3), which should not be the case in reality because the realistic accretion flow should be in vertical equilibrium (Chakrabarti 1996 and references therein); some results valid for isothermal accretion were directly applied to study the polytropic accretion in an ad-hoc manner (Chakrabarti 1996a), or some of the Newtonian approximations were not very convincingly combined with complete general relativistic equations (Chakrabarti 1996b) which does not strengthen their claim for a full general relativistic treatment of shock formation. Hence it is fair to say that although literature on general relativistic hydrodynamic BH accretion is well enriched by a number of important works (The following is an incomplete list of relevant papers on the subject: Novikov  \\& Thorne  1973, NT hereafter, Bardeen  \\& Petterson 1975, Abramowicz, Jaroszynski \\& Sikora 1978, Lu  1985,1986, Karas \\& Mucha 1993, Bj\\\"{o}rnsson 1995, Riffert  \\& Herold  1995, Ipser  1996, Pariev  1996, Peitz \\& Appl  1997,  Bao, Wiita \\& Hadrava  1998, Gammie \\& Popham  1998,  Gammie 1999), no well accepted complete general relativistic global shock solution exclusively obtained for a hydrodynamic accretion disc around a Schwarzschild BH has yet appeared in the literature.\\\\ Motivated by above mentioned limitations encountered by previous works in this field, the major aim of our work presented in this paper is to provide a generalized formalism which is expected to handle the formation of steady, standing Rankine- Hugoniot shock (RHS) in multi-transonic hydrodynamic BH accretion flow and to identify which region of parameter space (spanned by every important accretion parameter, namely, ${\\cal E}$, $\\lambda$ and $\\gamma$), will be responsible for such shock formation for {\\it all} available pseudo- Schawarzschild BH potentials. We would also like to compare the properties of multi transonic accretion in these BH potentials with complete general relativistic BH accretion as long as the issue of shock formation is concerned.\\\\ Hereafter, we will define the Schwarzschild radius $r_g$ as $$ r_g=\\frac{2G{M_{BH}}}{c^2} $$ (where  $M_{BH}$  is the mass of the black hole, $G$ is universal gravitational constant) so that the marginally bound circular orbit $r_b$ and the last stable circular orbit $r_s$ take the values $2r_g$ and $3r_g$ respectively for a typical Schwarzschild black hole. Also, total mechanical energy per unit mass on $r_s$ (sometimes called `efficiency' $e$) may be computed as $-0.057$ for this case. Also, we will use a simplified geometric unit throughout this paper where radial distance $r$ is scaled in units of $r_g$, radial dynamical velocity $u$ and polytropic sound speed $a$ of the flow is scaled in units of $c$ (the velocity of light in vacuum), mass $m$ is scaled in units of $M_{BH}$ and all other derived quantities would be scaled accordingly. Also, for simplicity,we will use $G=c=1$. In next section, we will briefly describe a few important features of the four different pseudo-Schwarzschild `effective' BH potentials used in this work. In \\S 3, we will show how we formulate and solve the equations governing multi-transonic BH accretion in these potentials which may have shocks. In \\S 4, we will study multi-transonic BH accretion using the full general relativistic frame work and will try to argue (in an indirect, but self consistant manner) which potential is expected to be the closest approximation of actual general relativistic solutions for which regions of parameter space spanned by ${\\cal E}$, $\\lambda$ and $\\gamma$, as long as one concentrates only on shocked flows. Finally in \\S 5 we will draw our conclusion by highlighting some of the possible important impacts of study of shock formation on related fields. \\section {Properties of four pseudo- Schwarzschild BH potentials} \\noindent Rigorous investigation of the complete general relativistic multi-transonic BH accretion disc structure is extremely complicated. At the same time it is understood that, as relativistic effects play an important role in the regions close to the accreting black hole (where most of the gravitational potential energy is released), purely Newtonian gravitational potential (in the form ${\\Phi}_{N}=-\\frac{1}{r}$) cannot be a realistic choice to describe transonic black hole accretion in general. To compromise between the ease of handling of a Newtonian description of gravity and the realistic situations described by complicated general relativistic calculations, a series of `modified' Newtonian potentials have been introduced to describe the general relativistic effects that are most important for accretion disk structure around Schwarzschild and Kerr black holes (see ABN for further discussion). Introduction of such potentials allows one to investigate the complicated physical processes taking place in disc accretion in a semi-Newtonian framework by avoiding pure general relativistic calculations so that most of the features of spacetime around a compact object are retained and some crucial properties of the analogous relativistic solutions of disc structure could be reproduced with high accuracy. Hence, those potentials might be designated as `pseudo-Kerr' or `pseudo- Schwarzschild' potentials, depending on whether they are used to mimic the space time around a rapidly rotating or non rotating/ slowly rotating (Kerr parameter $a\\sim0$) black holes respectively. Below we describe four such pseudo Schwarzschild potentials on which we will concentrate in this paper. It is important to note that as long as one is not interested in astrophysical processes extremely close (within $1-2~r_g$) to a black hole horizon, one may safely use the following BH potentials to study accretion on to a Schwarzschild black hole with the advantage that use of these potentials would simplify calculations by allowing one to use some basic features of flat geometry (additivity of energy or de-coupling of various energy components, i.e., thermal ($\\frac{a^2}{\\gamma-1}$), Kinetic ($\\frac{u^2}{2}$) or gravitational ($\\Phi$) etc., see subsequent discussions) which is not possible for calculations in a purely Schawarzschild metric (see \\S 4). Also, one can study more complex many body problems such as accretion from an ensemble of companions or overall efficiency  of accretion onto an ensemble of black holes in a galaxy  or  for studying numerical hydrodynamic accretion flows around a black hole etc. as simply as can be done in a Newtonian framework, but with far better accuracy. So we believe that a comparative study of multi-transonic accretion flow as well as shock formation using all these potentials might be quite useful in understanding some important features of various shock related astrophysical phenomena, at least until one can have a complete and self-consistent theory of complete general relativistic shock formation exclusively for a Schwarzschild BH.  However, one should be careful in using these potentials because none of these potentials discussed here are `exact' in a sense that they are not directly derivable from the Einstein equations. These potentials could only be used to obtain more accurate correction terms over and above the pure Newtonian results and any `radically' new results obtained using these potentials should be cross-checked very carefully with the exact general relativistic theory.\\\\ \\noindent Paczy\\'nski and Wiita (1980) proposed a pseudo-schwarzschild potential of the form $$ \\Phi_{1}=-\\frac{1}{2(r-1)} \\eqno{(1a)} $$ which accurately reproduces the positions of $r_s$ and $r_b$ and gives the value of efficiency to be $-0.0625$ which is in closest agreement with the value obtained in full general relativistic calculations. Also the Keplarian distribution of angular momentum obtained using this potential is exactly same as that obtained in pure Schwarzschild geometry. It is worth mentioning here that this potential was first introduced to study a thick accretion disc with super Eddington Luminosity. Also, it is interesting to note that although it had been thought of in terms of disc accretion, $\\Phi_1$ is spherically symmetric with a scale shift of $r_g$. \\\\ \\noindent To analyze the normal modes of accoustic oscillations within a thin accretion disc around a compact object (slowly rotating black hole or weakly magnetized neutron star), Nowak and Wagoner (1991) approximated some of the dominant relativistic effects of the accreting black hole (slowly rotating or nonrotating) via a modified Newtonian potential of the form $$ \\Phi_{2}=-\\frac{1}{2r}\\left[1-\\frac{3}{2r}+12{\\left(\\frac{1}{2r}\\right)}^2\\right ] \\eqno{(1b)} $$ $\\Phi_2$ has correct form of $r_s$ as in the Schwarzschild case but is unable to reproduce the value of $r_b$. This potential has the correct general relativistic value of the angular velocity $\\Omega_s$ at $r_s$. Also it reproduces the radial epicyclic frequency $\\kappa$ (for $r>r_s$) close to its value obtained from general relativistic calculations, and among all BH potentials, $\\Phi_2$ provides the best approximation for $\\Omega_s$ and $\\kappa$. However, this potential gives the value of efficiency as $-0.064$ which is larger than that produced by $\\Phi_1$, hence the disc spectrum computed using $\\Phi_2$ would be more luminous compared to a disc structure studied using $\\Phi_1$.\\\\ \\noindent Considering the fact that the free-fall acceleration plays a very crucial role in Newtonian gravity, ABN proposed two different BH  potentials to study disc accretion around a non-rotating black hole. The first potential proposed by them produces exactly the same value of the free-fall acceleration of a test particle at a given value of $r$ as is obtained for a test particle at rest with respect to the Schwarzschild reference frame, and is given by $$ \\Phi_{3}=-1+{\\left(1-\\frac{1}{r}\\right)}^{\\frac{1}{2}} \\eqno{(1c)} $$ The second one gives the value of the free fall acceleration that is equal to the value of the covariant component of the three dimensional free-fall acceleration vector of a test particle that is at rest in the Schwarzschild reference frame and is given by $$ \\Phi_{4}=\\frac{1}{2}ln{\\left(1-\\frac{1}{r}\\right)} \\eqno{(1d)} $$ Efficiencies produced by $\\Phi_3$ and $\\Phi_4$ are $-0.081$ and $-0.078$ respectively.The magnitude of efficiency produced by $\\Phi_3$ being maximum,calculation of disc structure using $\\Phi_3$ will give  the maximum amount of energy dissipation and the corresponding spectrum would be the most luminous one. Hereafter we will refer to all these four potentials by $\\Phi_i$ in general, where $\\left\\{i=1,2,3,4\\right\\}$ would correspond to $\\Phi_1$ (eqn. 1(a)), $\\Phi_2$ (eqn. 1(b)), $\\Phi_3$ (eqn. 1(c)) and $\\Phi_4$ (eqn. 1(d)) respectively. One should notice that while all other $\\Phi_i$s have singularity at $r=r_g$, only $\\Phi_2$ has a singularity at $r=0$. It can be shown that for $r>2r_g$, while $\\Phi_2$ is flatter compared to purely Newtonian potential $\\Phi_N$, all other $\\Phi_i$s are steeper to $\\Phi_N$.\\\\ \\noindent At any radial distance $r$ measured from the accretor, one can define the effective potential $\\Phi_i^{eff}(r)$ to be the summation of the gravitational potential and the centrifugal potential for matter accreting under the influence of $i$th pseudo potential. $\\Phi_i^{eff}(r)$ can be expressed as: $$ \\Phi_i^{eff}(r)=\\Phi_i(r)+\\frac{\\lambda^2(r)}{2r^2} \\eqno{(2a)} $$ where $\\lambda(r)$ is the non-constant distance dependent specific angular momentum of accreting material. One then easily shows that $\\lambda(r)$ may have an upper limit: $$ \\lambda^{up}_i(r)=r^{\\frac{3}{2}}\\sqrt{\\Phi^{'}_i(r)} \\eqno{(2b)} $$ where $\\Phi^{'}_i(r)$ represents the derivative of $\\Phi_i(r)$ with respect to $r$. For weakly viscous or inviscid flow, angular momentum can be taken as a constant parameter ($\\lambda$) and eqn. (2a) can be approximated as: $$ \\Phi_i^{eff}(r)=\\Phi_i(r)+\\frac{\\lambda^2}{2r^2} \\eqno{(2c)} $$ For general relativistic treatment of accretion, the effective potential can {\\it not} be decoupled in to its gravitational and centrifugal components. For a Schwarzschild metric of the following form: $$ ds^2=g_{{\\mu}{\\nu}}dx^{\\mu}dx^{\\nu} $$ $$ =-\\left(1-\\frac{1}{r}\\right)dt^2 +\\left(1-\\frac{1}{r}\\right)^{-1}dr^2 +r^2\\left(d{\\theta}^2+sin^2{\\theta}d{\\phi}^2\\right) $$ the world line of the accreting fluid is timelike, and the four velocity of the fluid satisfies the normalization condition: $$ u_{\\mu}u^{\\mu}=-1 $$ where $u^{\\mu}$($u_{\\mu}$) is the contra(co)-variant four velocity of the fluid. The angular velocity $\\Omega$ of the fluid can be computed as $$ \\Omega=\\frac{u^{\\phi}}{u^t}=-\\frac{{\\lambda}g_{tt}}{g_{{\\phi}{\\phi}}} =\\frac{{\\lambda}\\left(r-1\\right)}{2r^3} $$ where $\\lambda=-\\frac{u_{\\phi}}{u_t}$ is the specific angular momentum which is conserved for fluid dynamics as well as for particle dynamics for invicid flow. The  general relativistic effective potential $\\Phi^{eff}_{GR}(r)$ ({\\it excluding} the rest mass) experienced by the fluid accreting on to a Schwarzschild BH can be expressed as: $$ \\Phi^{eff}_{GR}(r)=r\\sqrt{\\frac{r-1}{r^3-{\\lambda}^2\\left(1+r\\right)}}-1 \\eqno{(2d)} $$ One can understand that the effective potentials in general relativity cannot be obtained by linearly combining its gravitational and rotational contributions because various energies in general relativity are combined together to produce non-linearly coupled new terms.\\\\ \\noindent In Fig 1, we plot $\\Phi_i^{eff}(r)$ (obtained from eq. (2c)) and $\\Phi^{eff}_{GR}(r)$ as a function of $r$ in logarithmic scale. The value of $\\lambda$ is taken to be 2 in units of $2GM/c$. $\\Phi^{eff}$ curves for different $\\Phi_i$ are marked exclusively in the figure and the curve marked by ${\\bf G^R}$ represents the variation of $\\Phi^{eff}_{GR}(r)$ with $r$. One can observe that $\\Phi^{eff}_1(r)$ is in excellent agreement with $\\Phi^{eff}_{GR}(r)$, only for a very small value of $r$ ($r{\\rightarrow}r_g$), $\\Phi^{eff}_1$ starts deviating from $\\Phi^{eff}_{GR}(r)$ and this deviation keeps increasing as matter approaches closer and closer to the event horizon. All other $\\Phi^{eff}_i(r)$s approaches to $\\Phi^{eff}_{GR}(r)$ at a radial distance (measured from the BH) considerably larger compared to the case for $\\Phi^{eff}_1(r)$. If one defines ${\\Delta}_i^{eff}(r)$ to be the measure of the deviation of $\\Phi^{eff}_i(r)$ with $\\Phi^{eff}_{GR}(r)$ at any point $r$, $$ {\\Delta}_i^{eff}(r)=\\Phi^{eff}_i(r)-\\Phi^{eff}_{GR}(r) $$ One observes that ${\\Delta}_i^{eff}(r)$ is always negative for $\\Phi^{eff}_1(r)$, but for other $\\Phi^{eff}_i(r)$, it normally remains positive for low values of $\\lambda$ but may become negative for a very high value of $\\lambda$. If ${{\\vert}}{\\Delta}_i^{eff}(r){{\\vert}}$ be the modules or the absolute value of ${\\Delta}_i^{eff}(r)$, one can also see that, although only for a very small range of radial distance very close to the event horizon, ${\\Delta}_3^{eff}(r)$ is maximum, for the whole range of distance scale while $\\Phi_1$ is the best approximation of general relativistic space time, $\\Phi_2$ is the worst approximation and $\\Phi_4$ and $\\Phi_3$ are the second and the third best approximation as long as the total effective potential experienced by the accreting fluid is concerned. It can be shown that ${{\\vert}}{\\Delta}_i^{eff}(r){{\\vert}}$ nonlinearly anti-correlates with $\\lambda$. The reason behind this is understandable. As $\\lambda$ decreases, rotational mass as well as its coupling term with gravitational mass decreases for general relativistic accretion material while for accretion in any $\\Phi_i$, centrifugal force becomes weak and gravity dominates; hence deviation from general relativistic case will be more prominent because general relativity is basically a manifestation of strong gravity close to the compact objects.\\\\ \\noindent From the figure it is clear that for $\\Phi^{eff}_{GR}(r)$ as well as for all $\\Phi^{eff}_i(r)$s, a peak appears close to the horizon. The height of these peaks may roughly be considered as the measure of the strength of the centrifugal barrier encountered by the accreting material for respective cases. The deleberate use of the word `roughly' instead of `exactly' is due to the fact that here we are dealing with fluid accretion, and unlike particle dynamics, the distance at which the strength of the centrifugal barrier is maximum, is located further away from the peak of the effective potential because here the total pressure contains the contribution due to fluid or `ram' pressure also. Naturally the peak height for  $\\Phi^{eff}_{GR}(r)$ as well as for $\\Phi^{eff}_i(r)$s increases with increase of $\\lambda$ and the location of this barrier moves away from the BH with higher values of angular momentum. If the specific angular momentum of accreting material lies between the marginally bound and marginally stable value, an accretion disc is formed. For invicid or weakly viscous flow, the higher will be the value of $\\lambda$, the higher will be the strength of the centrifugal barrier and the more will be the amount of radial velocity or the thermal energy that the accreting material must have to begin with so that it can be made to accrete on to the BH. In this connection it is important to observe from the figure that accretion under $\\Phi_1(r)$ will encounter a centrifugal barrier farthest away from the BH compared to other $\\Phi_i$s.  For accretion under all $\\Phi_i$s except $\\Phi_1$,the strength of centrifugal barrier at a particular distance will be more compared to its value for full general relativistic accretion. ",
        "conclusions": "\\noindent In this paper, we provide a generalized formalism which can formulate and solve the equations governing the advective, multi-transonic, hydrodynamic BH accretion in {\\it all} available pseudo-Schwarzschild potentials, which may contain steady, standing, Rankine-Hugoniot kind of shocks. We have also formulated and solved the equations governing multi-transonic, complete general relativistic BH accretion and wind in Schwarzschild metric and compared our pseudo-Schwarzschild solutions with the general relativistic one. The main conclusions of this paper are the following:\\\\ \\noindent (a) We observe that a significant region of parameter space (spanned by the conserved total specific energy ${\\cal E}$, the specific angular momentum $\\lambda$ and the polytropic index $\\gamma$ of the flow) allows shock formation for {\\it all} potentials, which leads to the strong conclusion that stable, standing RHS are inevitable ingredients in multi-transonic accretion disks around non-rotating BHs. The same kind of conclusion was drawn by previous works in this field  (see \\S 1) {\\it only} for ultra-relativistic accretion in Paczy\\'nski \\& Wiita (1980) potential, whereas we make this conclusion more general by incorporating {\\it all} available BH potentials to study BH  accretion  for {\\it all}  possible values of $\\gamma$.\\\\ \\noindent (b) As the shock forms at a particular radial distance, it is clear that self-similar solutions should {\\it not} be invoked while studying real physical BH accretion and related phenomena.\\\\ \\noindent (c) It is sometimes argued that a non-standing oscillating shock may modulate the disc spectrum in order to explain the dwarf novae outburst (Mausche, Raymond \\& Mattei 1995) or QPO (Hua, Kazanas \\& Titarchuk 1997). In this context, the region of parameter space, for which three sonic points are formed in accretion but still no steady, standing shock is found (see \\S 3), can be considered as  quite an important zone because $\\left[{\\cal P}_i\\right]\\in {\\cal G}_i\\left[{\\cal P}_i\\right]$ may provide the relevant parameters responsible for such physical processes.\\\\ \\noindent (d) As long as the shock formation in ultra-relativistic black hole accretion is concerned, the Paczy\\'nski \\& Wiita (1980) potential $\\Phi_1(r)$ is the {\\it only} pseudo potential which can mimic the solutions of general relativistic accretion disc around non-rotating BHs in a very efficient way. However, in the purely non-relativistic limit ($\\gamma{\\longrightarrow}5/3$), along with $\\Phi_1(r)$, another BH potential $\\Phi_4(r)$ proposed by ABN, is also observed to mimic the general relativistic solutions; at least for low energy - high angular momentum flows. However, it is interesting to note one important feature of the Paczy\\'nski \\& Wiita potential $\\Phi_1(r)$; like spherically symmetric accretion (see DS), for accretion disc also, $\\Phi_1(r)$ is observed to be in excellent agrement with solutions for ultra-relativistic flow in pure Schwarzschild metric, however, it starts loosing (albeight very slowly) its efficiency in mimicking full general relativistic solution with higher values of $\\gamma$, i.e., as the flow reaches its purely non-relativistic limits; although the exact reason behind this is not quite clear to us.\\\\ \\noindent Hot, dense and exo-entropic post-shock regions in advective accretion disks are used as a powerful tool in understanding the spectral properties of BH candidates (Shrader \\& Titarchuk 1998, and references therein) and in theoretically explaining a number of diverse phenomena, including millisecond variability in the X-ray emission from LMXBs and the generation mechanism for high frequency QPOs in general (Titarchuk, Lapidus \\& Muslimov 1998 and references therein), high energy emission from central engines of AGNs (Sivron, Caditz \\& Tsuruta 1996), formation of heavier elements in BH accretion discs via non-explosive nucleosynthesis (Mukhopadhyay \\& Chakrabarti 2000), formation and dynamics of accretion powered galactic and extragalactic jets, quiescent states of X-ray novae systems and outflow induced low luminosity of our galactic centre (Das 2001; Das \\& Chakrabarti 1999). A number of observational evidences are also present which are in close agreement with the theoretical predictions obtained from shocked accretion model (Rutledge et al. 1999; Muno, Morgan \\& Remillard 1999; Webb \\& Malkan 2000; Rao, Yadav \\& Paul 2000; Smith, Heindl \\& Swank 2001). Thus we believe that our present work may have far reaching consequences because of the following reasons: \\begin{enumerate} \\item Our generalized formalism assures that that our model is not just an artifact of a particular type of potential only and inclusion of every BH potential allows a substantially extended zone of parameter space allowing for the possibility of shock formation. \\item Of course there are possibilities that in future someone may come up with a pseudo-Schwarzschild potential better than $\\Phi_1(r)$, which will be the best approximation for complete general relativistic investigation of multi-transonic shocked flow. In such case, if one already formulates a generalized model for multi-transonic shocked accretion disc for any arbitrary $\\Phi(r)$, exactly what we have done in this paper, then that generalized model will be able to readily accommodate that new $\\Phi(r)$ without having any significant change in the fundamental structure of the formulation and solution scheme of the model and we need not have to worry about providing any new scheme exclusively valid only for that new potential, if any. \\item Even if someone can provide a completely satisfactory model for shock formation in full general relativistic (Schwarzschild) BH accretion, still the utility of this work may not be completely irrelevant. Rigorous investigation of some  of the shock related phenomena is extremely difficult (if not completely impossible) to study using full general relativistic framework. Hence one is expected to always rely on these pseudo-potentials because of the ease of handling them. For example, it was shown that (see \\S 4) the total energy of the general relativistic accretion flow can {\\it not} be decoupled into its constituent contributions, whereas for any kind of pseudo-potential (see \\S 2), all individual energy components are linear under addition. This provides enough freedom and ease to simply add any extra component in the expression for energy to introduce any new physics in the system (radiative forces or magnetic fields for example), which is certainly not possible while dealing with full general relativistic astrophysical flows around non-rotating BHs. \\end{enumerate} Thus, for above mentioned reasons, we believe that compared to all previous works based solely on ultra-relativistic accretion in $\\Phi_1$, our model  is better equipped for handling various shock related phenomena.\\\\ \\noindent It is noteworthy that the idea of shock formation in advective BH accretion is contested by some authors (Narayan, Kato \\& Honma 1997, and references therein). However, the fact that their claim against shock formation is,  perhaps, inappropriate for many reasons, has been shown (Molteni, Gerardi \\& Valenza 2001) from energy considerations. One can understand that the problem of not finding shocks lies in the fact that non-shock ADAF models are, perhaps, unable to produce multi-transonic flows because only one inner sonic point close to the BH is explored by such works. \\\\ \\noindent One can observe that flows characterized by $\\left[{\\cal P}_i\\right]\\in {\\cal F}_i\\left[{\\cal P}_i\\right]$ in our work may contain low intrinsic angular momentum for some cases (especially for purely non-relativistic flow in some of the BH potentials) However, such weakly rotating flows are expected to be allowed by nature for various real physical situations like detached binary systems fed by accretion from OB stellar winds (Illarionov \\& Sunyaev 1975; Liang \\& Nolan 1984), semi-detached low-mass non-magnetic binaries (Bisikalo et al.\\ 1998) and supermassive BHs fed by accretion from slowly rotating central stellar clusters (Illarionov 1988; Ho 1999 and references therein). \\\\ \\noindent Even twenty-eight years after the discovery of standard accretion disc theory (Shakura \\& Sunyaev 1973), exact modeling of viscous multi-transonic BH accretion, including proper heating and cooling mechanisms is still quite an arduous task, and we have not yet fully attempted this. However, our preliminary calculations show that the introduction of viscosity via a radius dependent power law distribution for angular momentum pushes the shock location closer to the BH; details of this work will be discussed elsewhere."
    },
    "0212/astro-ph0212433_arXiv.txt": {
        "abstract": "The dynamical equations describing the evolution of a self-gravitating fluid can be rewritten in the form of a Schr\\\"{o}dinger equation coupled to a Poisson equation determining the gravitational potential. This wave-mechanical representation allows an approach to cosmological gravitational instability that has numerous advantages over standard fluid-based methods. We explore the usefulness of the Schr\\\"{o}dinger approach by applying it to a number of simple examples of self-gravitating systems in the weakly non-linear regime. We show that consistent description of a cold self-gravitating fluid requires an extra ``quantum pressure'' term to be added to the usual Schr\\\"{o}dinger equation and we give examples of the effect of this term on the development of gravitational instability. We also show how the simple wave equation can be modified by the addition of a non-linear term to incorporate the effects of gas pressure described by a polytropic equation-of-state. ",
        "introduction": "The local Universe displays a rich hierarchical pattern of galaxy clustering that encompasses a large range of length scales, culminating in rich clusters and superclusters; the early Universe, however, was almost smooth, with only slight ripples seen in the cosmic microwave background radiation (e.g. Smoot et al. 1992). Models of the evolution of structure link these observations by appealing to the effect of gravitational instability, e.g. Lahav et al. (2002). Low-amplitude primordial perturbations of wave-like character within a largely homogeneous universe become amplified by their own self-gravity, first linearly which then, as their amplitudes build up, results in non-linear collapse. The linear theory of perturbation growth for cosmological density fluctuations is well-established, and is founded on tried-and-tested representation of matter on cosmological scales as a cold fluid (e.g. Peebles 1980). The non-linear regime is much more complicated and generally not amenable to analytic solution. Instead, numerical $N$--body simulations using particles rather than waves to trace cosmic structure, have led the way towards an understanding of strongly developed clustering. Analytic approaches based on particles, chiefly the Zel'dovich (1970) approximation, have also proved useful. Although these agree to an extent with fully numerical calculations (Coles, Melott \\& Shandarin 1993), they break down when particle trajectories cross to form singularities known as caustics.  In this paper we adopt an approach to cosmic structure that has advantages over both fluid and particle descriptions. Following Widrow \\& Kaiser (1993) and Coles (2002), we construct a formalism in which the dynamics of gravitational is couched in the language of quantum mechanics and governed by a Schr\\\"{o}dinger wave equation. The advantages of this approach will be made apparent as we explore its behaviour in simple examples which can be solved exactly and, within the Sch\\\"{o}dinger approach, to various levels of approximation. Our aim is to construct a useful tool for evolving density fluctuations into the mildly non-linear regime, beyond the breakdown of linear perturbation theory.  The importance of the search for analytic understanding of structure formation is often overlooked in the light of the tremendous advances that have been made recently in computational astrophysics. One motivation for a wider range of theoretical techniques is practical. High-resolution $N$-body experiments are CPU-intensive and it is difficult to find sufficient resources to use them to explore a large range of choices of initial fluctuations, different cosmological parameters, and so on. A neat analytic approximation can do this job more efficiently, at least to the point of narrowing down the field of contenders for fuller study. The other motivation is methodological: the complex ``cosmic web'' of large-scale structure requires {\\em explanation} rather than {\\em reproduction}, and explanation rests on using analytic approaches to identify the important physics as much as possible.  The layout of this paper is as follows. In Section 2 we run through the standard basics of gravitational instability in an expanding universe, including perturbation theory and the Zel'dovich approximation. We then, in Section 3, present the details of the Schr\\\"{o}dinger approach and derive a spherically-symmetric solution. In Section 4 we apply the method to examples of one-dimensional collapse showing how various approximations within the wave-mechanical framework behave when compared to the real data. We sum up in Section 5. ",
        "conclusions": "In this paper we have set out an approach to the formation of cosmic structure from evolving cosmological density fluctuations that relies on a transformation of the evolution equations into a Schr\\\"{o}dinger-Poisson system. One advantage of this new formalism was pointed out by Coles (2002): it yields a rather more convincing approach to understanding the origin of spatial intermittency in the galaxy distribution than that offered by Jones (1999). It also makes a connection in the underlying physics with other systems that display similar phenomena. In this paper we have argued further that it provides an analytical approach that offers benefits over standard approaches of both Lagrangian and Eulerian types. On the one hand, it always guarantees a positive density and on the other it does not produce singularities at shell-crossing events. It is clearly important to investigate both of these issues further. Szapudi \\& Kaiser (2003) have already completed a study of perturbation theory and its statistical ramifications within this approximation. Dealing fully with post-shell-crossing phenomena is beyond the scope of this paper as it requires a more sophisticated representation of the wave-function in velocity space; see Widrow \\& Kaiser (1993) for discussion.  In this paper we showed how this approach works in practice by applying it to a static spherically-symmetric case and to some examples of one-dimensional collapse. In doing this we looked at a ``free-particle'' solution of the system that resembles the Zel'dovich approximation in some respects. Incorporating a fixed gravitational potential yields a solution that has some similarities to the ``frozen-potential'' approximation (Matarrese et al. 1992). We think the Schr\\\"{o}dinger formalism provides a useful tool for modelling fluctuations in the mildly non-linear regime. In future work we will look at its use in modelling the intergalactic medium in order to understand properties of Lyman-$\\alpha$ absorption systems."
    },
    "0212/astro-ph0212575_arXiv.txt": {
        "abstract": "{ Observations of dark matter dominated dwarf and low surface brightness disk galaxies favor density profiles with a flat-density core, while cold dark matter (CDM) N-body simulations form halos with central cusps, instead. This apparent discrepancy has motivated a re-examination of the microscopic nature of the dark matter in order to explain the observed halo profiles, including the suggestion that CDM has a non-gravitational self-interaction. We study the formation and evolution of self-interacting dark matter (SIDM) halos. We find analytical, fully cosmological similarity solutions for their dynamics, which take proper account of the collisional interaction of SIDM particles, based on a fluid approximation derived from the Boltzmann equation. The SIDM particles scatter each other elastically, which results in an effective thermal conductivity that heats the halo core and flattens its density profile. These similarity solutions are relevant to galactic and cluster halo formation in the CDM model. We assume that the local density maximum which serves as the progenitor of the halo has an initial mass profile \\( \\delta M/M\\propto M^{-\\varepsilon } \\), as in the familiar secondary infall model. If \\( \\varepsilon =1/6 \\), SIDM halos will evolve self-similarly, with a cold, supersonic infall which is terminated by a strong accretion shock. Different solutions arise for different values of the dimensionless collisionality parameter, \\( Q\\equiv\\sigma \\rho_{b}r_{s} \\), where \\( \\sigma  \\) is the SIDM particle scattering cross section per unit mass, \\( \\rho _{b} \\) is the cosmic mean density, and \\( r_{s} \\) is the shock radius. For all these solutions, a flat-density, isothermal core is present which grows in size as a fixed fraction of \\( r_{s} \\). We find two different regimes for these solutions: 1) for \\( Q<Q_{th}(\\simeq 7.35\\times 10^{-4}) \\), the core density decreases and core size increases as \\( Q \\) increases; 2) for \\( Q>Q_{th} \\), the core density increases and core size decreases as \\( Q \\) increases. Our similarity solutions are in good agreement with previous results of N-body simulation of SIDM halos, which correspond to the low-\\( Q \\) regime, for which SIDM halo profiles match the observed galactic rotation curves if \\( Q\\sim[8.4\\times 10^{-4} - 4.9 \\times 10^{-2}]Q_{th}\\), or \\( \\sigma \\sim [0.56 - 5.6]\\, cm^{2} g^{-1} \\). These similarity solutions also show that, as \\( Q \\rightarrow \\infty \\), the central density acquires a singular profile, in agreement with some earlier simulation results which approximated the effects of SIDM collisionality by considering an ordinary fluid without conductivity, i.e. the limit of mean free path \\( \\lambda_{mfp} \\rightarrow 0 \\). The intermediate regime where \\( Q\\sim [18.6-231]Q_{th} \\) or \\( \\sigma \\sim [1.2 \\times 10^{4} - 2.7 \\times 10^{4}]\\,cm^{2}g^{-1} \\), for which we find flat-density cores comparable to those of the low-\\( Q \\) solutions preferred to make SIDM halos match halo observations, has not previously been identified. Further study of this regime is warranted. } ",
        "introduction": "The cold dark matter (CDM) model provides a successful framework for understanding the formation and evolution of structure in the universe. According to this model, structure forms out of primordial density fluctuations by hierarchical clustering; small objects form first, and then merge to make bigger objects. However, this CDM model has several problems.  There has been a recent concern about the possible discrepancy in the inner density profile of dark-matter dominated cosmological halos between N-body simulations and observations. CDM halos in N-body simulations have a density cusp (\\( \\rho \\propto r^{\\beta } \\) where \\( \\beta  \\) ranges from -1 (Navarro, Frenk \\& White 1997; hereafter ``NFW'') to -1.5 (Moore et al. 1999; hereafter ``Moore profile''), while observations of dwarf and LSB galaxies, which are believed to be dark-matter dominated, indicate flat-density (soft) cores.  Spergel \\& Steinhardt (2000) suggested that the purely collisionless nature of CDM be replaced by self-interacting dark matter (SIDM) to resolve this discrepancy. The non-gravitational, microscopic interaction (e.g. elastic scattering) of SIDM particles is postulated to be strong enough to produce a soft core by heat conduction. Cosmological N-body simulations which incorporate a finite scattering cross-section  \\( \\sigma  \\) for the SIDM particles show that this scheme successfully produces soft cores in halos (e.g. Dav\\'{e} et al. 2001; Yoshida et al. 2000b).  Previous analytical study (Balberg, Shapiro \\& Inagaki 2002) and estimates based on N-body simulations (Burkert 2000; Kochanek \\& White 2000) indicate that SIDM cores will become unstable by undergoing gravothermal catastrophe in a Hubble time. This potential instability of SIDM cores means that, if the matter content of the universe is really SIDM particles, we are living in a special epoch where SIDM halos still have soft cores. However, these analyses are based on isolated halos. In the context of the hierarchical clustering scenario, cosmological infall is inevitable. We are led to pose the following question: Can cosmological infall delay the core collapse of SIDM halos?  In this paper we study the formation and evolution of SIDM halos with a proper treatment of cosmological infall. We find that, for a range of infall rates, the collapse (i.e. gravothermal catastrophe) of SIDM cores can be completely inhibited. We also find a new range of values of the scattering cross section which can produce soft cores, separate from and in addition to the range previously identified by N-body simulations. Toward this end, we find a self-similar solution to describe the evolution of SIDM halos for arbitrary degree of collisionality which allows a detailed analytic study. ",
        "conclusions": "We have found similarity solutions which allow us to describe the dynamical origin and evolution of SIDM halos in a fully cosmological context, analytically, for the first time. Our solutions are based upon fluid-like conservation equations which can be derived by taking moments of the Boltzmann equation, assuming that the particle velocity distribution inside virialized halos is approximately isotropic, as CDM N-body simulations suggest. Our solutions confirm and explain the results of N-body simulations which have attempted to incorporate the collisional effects of SIDM numerically within the CDM model and allow us to extend and generalize those results to a much wider range of parameters and scales than has so far been simulated. Along the way, we have also found a similarity solution for the adiabatic case of standard CDM (i.e. with no SIDM self-interaction) which serves to derive analytically the halo profile of CDM halos found previously by N-body simulation (i.e. intermediate between NFW and Moore profiles), as a product of self-similar cosmological infall.  As seen in our similarity solutions, cosmological infall can affect the dynamics of virialization significantly. In the context of hierarchical clustering, in fact, the collapse of halo cores previously predicted for \\(isolated\\) SIDM halos is entirely prevented by cosmological infall, according to these similarity solutions. For realistic mass accretion histories in a CDM universe, therefore, SIDM core collapse will be delayed until this infall becomes negligible. Accordingly, previous analyses (Burkert 2000; Balberg, Shapiro \\& Inagaki 2002) which predict core collapse in a Hubble time should be re-examined.  By considering the full range of collisionality from \\(Q=0\\) to \\(Q=\\infty\\), we have discovered that the low-\\(Q\\) regime which has been found previously by N-body simulation to produce observationally acceptable soft cores is not unique. We predict that a cross-section of the order of \\( 10^{4}\\,cm^{2}g^{-1} \\) can also produce acceptable soft cores, in addition to the regime of smaller cross-section previously identified. Therefore, we suggest that cosmological N-body simulations be performed which incorporate SIDM particles with \\( \\sigma =[5\\times 10^{3}-5\\times 10^{4}]\\,cm^{2}g^{-1} \\) to explore this possibility further.  In reality, the mass accretion rates of CDM halos in N-body simulations are found to depart from the self-similar values over time, declining at late times rather than growing (Wechsler et al. 2002). We have accordingly begun to study the evolution of SIDM halos with a more realistic infall, using a 1-D, spherical hydrodynamics code, and will describe those results elsewhere (Ahn and Shapiro, in preparation).  \\vskip 0.4cm This work was partially supported by NASA ATP Grant NAG5-10825 and Texas Advanced Research Program Grant 3658-0624-1999."
    },
    "0212/astro-ph0212096_arXiv.txt": {
        "abstract": "{We present a new method to analyse and reduce chemical networks and apply this technique to the chemistry in molecular clouds. Using the technique, we investigated the possibility of reducing the number of chemical reactions and species in the UMIST\\,95 database simultaneously. In addition, we did the same reduction but with the ``objective technique'' in order to compare both methods. We found that it is possible to compute the abundance of carbon monoxide and fractional ionisation accurately with significantly reduced chemical networks in the case of pure gas-phase chemistry. For gas-grain chemistry involving surface reactions reduction is not worthwhile. Compared to the ``objective technique'' our reduction method is more effective but more time-consuming as well. ",
        "introduction": "Despite the significant growth of computer power over the last years, the modelling of the chemical evolution of protostellar cores and protoplanetary discs is still a challenging computational task. This is even more true if one aims at the simultaneous self-consistent modelling of the dynamical and chemical evolution of these systems.  The mathematical formulation of the chemical processes in space is realised by differential equations describing chemical kinetics. They contain terms for all chemical reactions expected to proceed efficiently in the interstellar medium (ISM). Along with abundances of all involved species, these terms depend on the rate coefficients that are summarised in special databases. The relevant astrochemical rate files, like the UMIST rate file or Ohio New Standard Database, comprise hundreds of species and thousands of reactions (Millar et al.~\\cite{UMIST90,UMIST95}; Le Teuff et al.~\\cite{UMIST99}; Terzieva \\& Herbst~\\cite{TH98}; Aikawa \\& Herbst~\\cite{AH99}). Of course, one can only wonder if these databases contain all the information about reactions needed to accurately reproduce observed abundances of interstellar molecules. The opposite question to ask is if all this information is always needed.  In many astronomically interesting situations it is enough to follow the evolution of a limited number of species only. Intuitively, it seems to be apparent that to predict an abundance of N$_2$H$^+$, for example, one can afford computing C$_{10}$ abundance with less accuracy or even ignore this molecule altogether. In most cases, though, modelling of the chemical evolution of a parcel of the interstellar gas under fixed conditions or in a limited predefined range of conditions is not very time-consuming, so one probably will not be eager to sacrifice any accuracy even for the sake of order of magnitude gain in computational speed. However, in the case of dynamical modelling with changing density, UV intensity, temperature etc., a reduced chemical network may help to distinguish between feasible and non-feasible problems or at least to ease the computational burden if parallel processors are in use. One possible application for reduced chemical networks would be the modelling of the evolution of magnetised protostellar clouds or protoplanetary accretion discs, when it is necessary to compute the fractional ionisation self-consistently. Another example of dynamically important species is carbon monoxide (CO) which under certain conditions can be the most important molecular coolant (see, for example, Glassgold \\& Langer~\\cite{GL73}).  Several authors have already used reduced chemical networks, based on the above arguments, however, without presenting accurate mathematical considerations (e.g., Gerola \\& Glassgold~\\cite{GG80}; Howe, Hartquist, \\& Williams \\cite{howe}; El-Nawawy, Howe, \\& Millar~\\cite{El97}). Recently, Ruffle et al.~(\\cite{Rea02}) (hereafter RRPHH) and Rae et al.~(\\cite{Rae02}) (hereafter RBHPR) made a first attempt to reduce astrochemical networks by so-called ``objective reduction techniques'', developed in combustion chemistry. They have shown that it is possible to quantify the reduction task by introducing a certain criterion that would allow in each particular case to indicate, abundances of which species must be followed in order to have a reasonable estimate for an abundance of a species under investigation. For static regions with physical conditions typical of diffuse clouds, they isolated the species set that contains tens instead of hundreds species. This reduction only introduces a factor of two uncertainty in predicted CO abundance. The same task has been solved for the fractional ionisation as well.  We perform a similar analysis of the UMIST\\,95 chemical reaction database in order to find if it is possible to reduce the number of species and/or reactions in the chemical networks used to calculate ionisation degree and carbon monoxide abundance under conditions typical of molecular clouds. Unlike RRPHH and RBHPR, we partly focus on the medium that is more advanced toward the formation of a star. This situation is more complicated from the chemical point of view. First, pre-stellar clouds are often shielded from the interstellar UV radiation field. Dissociating photons, being an additional chemical factor, in effect simplify the chemical network, as they prevent many complex species from being formed. Reactions in grain surfaces are another factor that cannot be ignored in dense clouds.  The paper is organised as follows. First, we describe our chemical model in Sect.~2. The species-based and reaction-based reduction methods are outlined in Sect.~3. Results of reductions obtained with both methods are presented in Sect.~4 (ionisation degree) and Sect.~5 (CO). These results are discussed in Sect.~6. A conclusion follows. ",
        "conclusions": "We developed a robust method to reduce simultaneously the number of species and reactions in chemical networks. The straightforward application for such reduced chemical networks would be the modelling of the evolution of magnetised protostellar clouds or protoplanetary discs, when it is necessary to compute the fractional ionisation self-consistently with dynamical processes.  Investigating the applicability of this reduction approach, we considered the conditions of diffuse and dense molecular clouds and utilised the UMIST\\,95 database of chemical reaction rates. For the sake of comparison, we performed the same reduction with the objective reduction technique. The new reaction-based way of reduction proved to be more efficient and accurate but more time-consuming as well.  It is found that the number of species and reactions that are needed to follow the evolution of the CO abundance or the ionisation degree can be greatly decreased in the case of pure gas-phase chemistry, implying computational speed gains of more than a few hundred. For instance, to follow accurately (with $<15\\%$ uncertainty) the evolution of the CO abundance in a dense cloud, it is enough to retain only 8 species involved in 9 reactions in the chemical network. However, as soon as the gas-grain interactions and surface chemistry are taken into account, reduction is modest and corresponding acceleration factors are around 10."
    },
    "0212/astro-ph0212269_arXiv.txt": {
        "abstract": " ",
        "introduction": "The scientific interest in TeV blazars is manifold and ranges from classification studies of the AGN (active galactic nucleus) phenomenon and detailed studies of Doppler-boosted relativistic jets, to particle acceleration and $\\gamma$-ray emission models near the central massive object. In addition, $\\gamma$-ray beams traversing intergalatic space can be used to constrain and measure the diffuse extragalactic background light (EBL) through $\\gamma\\gamma \\rightarrow e^{+}e^{-}$ absorption. In this paper we present recent results relevant to those topics.      \\bigskip ",
        "conclusions": ""
    },
    "0212/hep-ph0212112_arXiv.txt": {
        "abstract": "The so-called curvaton mechanism --a way to convert isocurvature perturbations into adiabatic ones-- is investigated both analytically and numerically in a pre-big bang scenario where the r\\^ole of the curvaton is played by a sufficiently massive Kalb--Ramond axion of superstring theory. When combined with observations of CMBR anisotropies at large and moderate angular scales, the present analysis allows us to constrain quite considerably the parameter space of the model: in particular, the initial displacement of the axion from the minimum of its potential and  the rate of evolution of the compactification volume during pre-big bang inflation. The combination of theoretical and experimental constraints favours a slightly blue spectrum of scalar perturbations, and/or a value of the string scale in the vicinity of the SUSY-GUT scale. ",
        "introduction": "\\label{Sec1} At present, the largest-scale temperature fluctuations of the Cosmic Microwave Background radiation (CMBR) are consistent with a (quasi) scale-invariant spectrum of  Gaussian primordial curvature fluctuations \\cite{cobe,cob2,silk}. The analysis of the first acoustic oscillations occurring on shorter angular scales adds the information that such curvature fluctuations should  be predominantly adiabatic \\cite{boom,dasi,maxima,archeops}. Although sufficiently small amounts of non-Gaussianity and/or isocurvature perturbations are not excluded, the above-mentioned observational features represent an important constraint for any scenario trying to model the initial stages of our Universe. In a previous Letter \\cite{us} we  tried to confront the  pre-big bang  scenario \\cite{g1,lwc,g2} with these constraints. The present paper contains a full description of that analysis and completes it.  Let us recall that, during the pre-big bang phase, the quantum fluctuations of all the light modes present in the low energy effective action are parametrically  amplified. Nonetheless, sizeable large-scale adiabatic fluctuations are not easily produced from the initial vacuum through the usual mechanism of parametric amplification. In particular, both tensor and scalar-metric fluctuations are amplified with  very steep spectra \\cite{GV94,bmggv}, resulting in  adiabatic modes which are far too small to explain the observed level of large-scale CMBR anisotropies \\cite{smoo1}.  However, not all the primordial spectra of pre-big bang cosmology are blue. For instance, in a pure gravi-dilaton background, the pseudo-scalar supersymmetric  partner of the dilaton in the dimensionally-reduced string effective action, the so-called Kalb--Ramond axion, emerges from the pre-big bang phase with a fluctuation spectrum whose tilt depends on the rate of change of  the compactification volume \\cite{cew,clw}. Depending on this, the axion tilt can be negative (red spectrum), positive (blue spectrum), or zero (scale-invariant spectrum). However, since the homogeneous (background) component of the axion is trivial, such a spectrum does not affect directly the metric, hence no curvature perturbation is generated at this primordial level. In other words, the axion perturbations are entropic, isocurvature perturbations. This feature persists, unfortunately, under $S$-duality transformations \\cite{cew,clw}. If such an axion is massless, or at least light enough not to have decayed yet , the induced CMBR fluctuations on large scales can fit the COBE normalization \\cite{dgs1,dgs2,dgs3,mas1}, but, being neither adiabatic nor  Gaussian \\cite{mvdv}, they are not able to fit the observed  structure of the first few  acoustic peaks.  A possible way out of this problem \\cite{us,lwc,es,lw,Moroi} is  offered by the alternative scenario of a massive axion, initially displaced from the minimum of its non-perturbative potential.  In that case axion perturbations couple to scalar metric perturbations through the non-vanishing axion's VEV. Eventually, the axion relaxes toward the minimum  of the potential and then, if heavy enough, decays prior to nucleosynthesis. During the relaxation process  the dominant source of energy undergoes a drastic change: it consists of the radiation produced at the  end of the pre-big bang evolution, and later  becomes the pressureless fluid  corresponding to the damped coherent oscillations of the axion. This non-trivial evolution results in a non-adiabatic pressure perturbation which, in  turn, is well known \\cite{ks,mfb} to induce  curvature perturbations on constant energy (or comoving) hypersurfaces even on superhorizon scales.  The interplay of such different sources of inhomogeneity, throughout the different stages of the background evolution, eventually determines the spectral amplitude of scalar curvature perturbations right after matter-radiation equality, when  all the scales of interest for the CMBR data are still outside the horizon. This conversion of isocurvature into adiabatic perturbations, originally suggested in a  different context by Mollerach \\cite{mol}, also applies to more general cases \\cite{lw,Moroi}.   Depending upon the initial value $\\sigma_{\\rm i}$ of the  Kalb--Ramond background,  different post-big bang histories are possible. If $\\bar{\\sigma} < \\sigma_{\\rm i} \\ll 1$ in Planck units (see below, Eq. (\\ref{c1}), for the definition of $\\bar{\\sigma}$), the axion  oscillates for a long time  before  becoming dominant and eventually decays. For $  \\sigma_{\\rm i} < \\bar{\\sigma}$ it may never fully dominate the energy density before decaying. If, instead, $\\sg_{\\rm i} \\gg 1$  the axion will dominate before oscillating  and a  slow-roll (low-scale) inflationary  phase could take place in that epoch. As we shall see, of all these possibilities CMBR observations seem to  favour the ``natural\" one, $\\sigma_{\\rm i} \\sim 1$. In any case, even if different post-big bang histories will lead to different spectral amplitudes of the Bardeen potential, adiabatic scalar metric perturbations will always be present at some level outside the horizon, prior to decoupling.  The purpose of the present paper is to report on the calculation of the spectral amplitude of the induced adiabatic metric perturbations, and on the comparison of the predictions of the pre-big bang scenario with the observations coming from the physics  of the CMBR anisotropies. In order to achieve this goal it is mandatory to have a good understanding both of the axion relaxation mechanism and of the evolution of the inhomogeneities. Hence analytical results will be supported with numerical examples and vice versa. We will present, in particular,  a full derivation of the results for the final adiabatic spectrum of the Bardeen potential (some of these results have been summarized already in \\cite{us}).  The paper is organized as follows. In Section II the basic equations describing the post-big bang evolution of the inhomogeneities and of the background geometry will be introduced. In Section III the physics of the  different post-big bang histories will be analysed. In Section IV the evolution of the background and of its perturbations will be discussed for the case in which the amplitude of the initial axion background is smaller than $1$ in Planck units, $\\sg_{\\rm i}<1$. In Section V we will discuss the evolution of the system in the complementary case $\\sg_{\\rm i}>1$. Section VI is devoted to the phenomenological implications of  the large-scale adiabatic perturbations produced through the relaxation of the axionic background. The obtained results will be compared with observations. Constraints on the pre-big bang parameters will be derived. Section VII contains our concluding remarks while, in the Appendix, a self-contained derivation of the axionic spectra produced by the pre-big bang evolution has been included.  \\renewcommand{\\theequation}{2.\\arabic{equation}} \\setcounter{equation}{0} ",
        "conclusions": "In the present paper the possible conversion of isocurvature, primordial axionic fluctuations into adiabatic, large-scale metric perturbations has been discussed in the context of the pre-big bang scenario.  Depending upon the specific relaxation of the axionic background toward the minimum of the potential, a constant (and large enough) mode  in the Bardeen potential can be generated, for scales  that are still outside the horizon right after matter-radiation equality.  After analysing the dynamics of the background and of its fluctuations, the final  amplitude and spectrum of the Bardeen potential has been related to the initial axion spectrum directly arising from the vacuum fluctuations amplified during the pre-big bang epoch. Our goal has been to include, with reasonable accuracy, the details of the post-big bang evolution, in  such a way that the pre-big bang parameters could be directly constrained by the COBE normalization, and by the analysis of the Doppler-peak structure.  All the theoretical uncertainty reflects in our lack of knowledge of $H_1$ which determines the end point of the primordal axion spectrum.  The main conclusion of this work is that a phenomenologically appealing spectrum of adiabatic scalar perturbations can naturally emerge from the simplest pre-big bang scenario through a conversion of the initial isocurvature perturbations of the Kalb-Ramond axion. Since, at the large scales tested by CMBR experiments, the above conversion preserves the scale-dependence of the original spectrum, it is important for the latter to be quasi-scale-invariant at large scales. This can be achieved, for instance, if the very early stages of pre-big bang cosmology at weak coupling involve a symmetric evolution of all $9$ spatial dimensions (modulo $T$-duality). Since the constant mode of the curvature fluctuations leads to adiabatic initial conditions for the fluid evolution after matter-radiation equality, the location of the Doppler peak is correctly reproduced.  On the other hand, the absolute normalization of fluctuations at large scales (say those relevant for COBE) depend on several details of the model. Indeed, the axion spectrum is naturally normalized at its end-point, given by our parameter $H_1$. If one takes, naively, $H_1 \\sim M_{\\rm s} \\sim 10^{17} ~{\\rm GeV}$ and assumes a flat spectrum one finds values of $\\Delta T/T$ that are a couple of orders of magnitude too large when compared with COBE's data. However, one can think of many (individual or combined) effects that can bring down our normalization to agree with the data, e.g. \\begin{itemize} \\item A slight (blue) tilt to the spectrum; \\item A blue spectrum just at high frequency (i.e. for scales that exit late, during the strongly coupled regime); \\item A lower  $H_1/ M_{\\rm s}$  ratio; \\item A  lower $M_{\\rm s}/ M_{\\rm P}$ ratio. \\end{itemize}  In the near future we hope to extend the present discussion to forthcoming CMBR anisotropy data at even smaller angular scales. It would be interesting to see if a combined analysis of the experimental data may give further useful hints on the parameter space of the scenario explored in the present investigation."
    },
    "0212/astro-ph0212082_arXiv.txt": {
        "abstract": "X-ray studies of high-redshift ($z>4$) active galaxies have advanced substantially over the past few years, largely due to results from the new generation of X-ray observatories. As of this writing X-ray emission has been detected from nearly 60 high-redshift active galaxies. This paper reviews the observational results and their implications for models of the first massive black holes, and it discusses future prospects for the field. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/astro-ph0212561_arXiv.txt": {
        "abstract": "The thermalized energy from the radioactive decays of $^{56}$Ni and $^{57}$Ni and their daughter nuclides power the light curves of supernovae near maximum light. The bolometric light curve gives us a fundamental understanding of the energy evolution of a supernova explosion and the amount of radioactive nuclides produced. In this review, I will discuss the bolometric evolution of the Type IIp supernovae SN1987A, and the general class of bolometric light curves of Type Ia thermonuclear explosions. ",
        "introduction": "We have been monitoring the photometric properties of the Type IIp SN1987A in the LMC for more than 15 years. Over the first 5 years the optical photometry could be measured from the ground in the typical 1'' seeing at our observatories. In the near infrared bands of $JHK$ where the seeing is better and the crowding stars (which are typically bluer than the SN) are less of a problem, we could monitor the SN for up to 10 years. Once the supernova faded below the integrated magnitudes of the inner ring which lies at about 1'' from the supernova debris, $HST$ or ground-based AO imagery is needed to isolate the debris and inner ring evolution. In Figure \\ref{fig01}, I show the structure of the inner ring region.  \\begin{figure}[h] \\begin{center} \\includegraphics[angle=0,width=0.5\\textwidth]{nsuntzeff_fig01.eps} \\end{center} \\caption[nsuntzeff_fig01.eps]{An $HST$ image of region near the inner ring of SN1987A. Region 1 isolates the expanding debris. Region 2 represents the shocked region between the debris and the inner ring. Region 3 represents the inner ring} \\label{fig01} \\end{figure}  From our $HST$ images, we have measured the evolution of the brightness of the three regions shown in Fig.~\\ref{fig01} with respect to an averaged background region outside of Region 3. For the earlier ground-based data where we can only measure the light from all three regions, we have subtracted off the extrapolated flux from Regions 2 and 3 to derive estimated fluxes for just the debris region. In Fig.~\\ref{fig02} we plot the optical light curve for 15 years of evolution. With these corrections, the optical and $HK$ photometry from the ground and $HST$ flight system equivalents merge well together. The near-infrared photometry in $J/F110W$ shows a large discontinuity presumably due to the radical differences in the two filter systems. The photometry of non-stellar SEDs can show large systematic errors due to the differences in the sensitivity functions of the atmosphere/telescope/filter/detector system. In early photometry of SN1987A, we noted difference of up to 0.4\\,mag in $I$ photometry due to the differences in the facility $I$ filters \\cite{Sun_etal99}. In Table \\ref{tab01}, I list the latest photometry for SN1987A  \\begin{table} \\caption{Optical Photometry of SN1987A as of May 2002} \\begin{center} \\renewcommand{\\arraystretch}{1.0} \\setlength\\tabcolsep{5pt} \\begin{tabular}{lccccc} \\hline\\noalign{\\smallskip} & $U$ & $B$ & $V$ &$R$ &$I$ \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} debris & 20.91(32) & 20.88(12) & 21.28(07) & 20.96(22) & 20.46(04)\\\\ decay rate (mag y$^{-1}$) & 0.23  & 0.20  & 0.24  & 0.35  & 0.22 \\\\ \\hline \\end{tabular} \\end{center} \\label{tab01} \\end{table}  \\begin{table} \\caption{Near-Infrared Photometry of SN1987A} \\begin{center} \\renewcommand{\\arraystretch}{1.0} \\setlength\\tabcolsep{5pt} \\begin{tabular}{cccc} \\hline\\noalign{\\smallskip} days since explosion& $J$ & $H$ & $K$ \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 4151 & 19.81(09) & 18.29(06) & 18.40(02)\\\\ 5428 & 20.46(15) & ... & ... \\\\ \\hline \\end{tabular} \\end{center} \\label{tab02} \\end{table}  The leveling off of the light curves after year 5 is due to two causes: the longer decay times for the remaining radioactive nuclides and the freeze-out of the cooling. By year five, the main radioactive energy sources of $^{56}$Ni and its daughter nucleus $^{56}$Co with e-folding times of 8.8d and 111.3d have decayed away. The nuclides $^{57}$Co (the daughter nuclide of $^{57}$Ni) and $^{44}$Ti with e-folding times of 390d and 87y are left as the main energy input, along with the energy input from a possible (but as yet unseen) pulsar. However, the light curves at this phase no longer represent the prompt thermalization of the input radioactive energy. As shown in \\cite{Fra_Koz93,Fra_Koz01}, the recombination and cooling time scales become larger than the expansion time scale, and the debris nebula is not able to cool at the rate of radioactive energy input.   \\begin{figure}[h] \\begin{center} \\includegraphics[angle=0,width=0.95\\textwidth]{nsuntzeff_fig02.eps} \\end{center} \\caption[nsuntzeff_fig02.eps]{Optical light curves (left panel) and near-infrared light curves (right panel) for the debris (Region 1) in SN1987A. The data have been shifted by the magnitudes listed in the legend. The data are a combination of ground-based $UBVRIJHK$ and the equivalent $HST$ flight system magnitudes.  } \\label{fig02} \\end{figure}  The ultraviolet, optical, and infrared photometry can be integrated to the thermalized ``uvoir'' energy flux to derive the masses of the synthesized radioactive nuclides. The early time data (less than 1000d) with prompt thermalization and cooling can be simply fit to radioactive decay models as in \\cite{Sun_Bou90,Bou_etal91,Sun_etal92}. The late-time bolometric light curves, when fit with models including the freeze-out, can be used to measure the amount of $^{57}$Ni and $^{44}$Ti. Fransson and Kozma (\\cite{Fra_Koz01}) present the latest values for the amounts of these nuclides where they found the following masses: M($^{57}$Ni,$^{57}$Ni,$^{44}$Ti)=(0.069,0.003,0.0001)M$_\\odot$. Lundqvist et al.~(\\cite{Lun_etal01}) find upper limits on the mass of $^{44}$Ti consistent with these numbers based on the non-detection of the 25m$\\mu$ lines of [FeI] and [FeII] from ISO/SWS. They caution, however, that if dust cooling is important, the limits on the mass of $^{44}$Ti could be significantly higher. Only direct detection of the 1.157 MeV line of $^{44}$Ti will resolve the ambiguity in the models.  In Fig.~\\ref{fig03} I show the late-time photometric data for the ring and the inter-ring region. The ring region suddenly began to brighten at a rate of $\\sim -0.24$ mag y$^{-1}$ in $UBVI$ and $\\sim -0.12$ mag y$^{-1}$ in $R$ since year 13. The brightening of the ring was predicted in \\cite{Lou_McC91,Che_Dwa95,Bor_etal97}. The ring was expected to brighten by 7.5mag to about $V=11$ as the blast wave from the debris struck the inner ring starting around year 16-20. However, in year 10, Pun et al.~(\\cite{Pun_etal97}) discovered a single spot brightening, evidently due to an inward protrusion of the inner ring. Our observations here show that the general rapid brightening of the ring began later, around year 13.  The inter-ring region has accelerated its brightening since year 13 and is now brightening at a rate of $\\sim -0.12$ mag y$^{-1}$ in $UBVI$. The inter-ring light is presumably due to the the emission from the reverse shock which is located at $\\sim75$\\% of the radius of the inner boundary of the inner ring.  This emission was predicted by \\cite{Bor_etal97} and discovered by \\cite{Son_etal98} with STIS data from $HST$. Refined models are presented in \\cite{Mic_etal98a,Mic_etal98b}. The reverse shock formed in the equatorial plane as neutral hydrogen atoms streaming from the debris at 15000 km s$^{-1}$ hit the ionized region at the inner surface of the ring. There is also diffuse light present in the STIS data which may be due to excitation by non-thermal particles accelerated by the shock. \\begin{figure}[h] \\begin{center} \\includegraphics[angle=0,width=0.8\\textwidth]{nsuntzeff_fig03.eps} \\end{center} \\caption[nsuntzeff_fig03.eps]{$UBVRI$ optical light curves for the ring Region 3 (left panel) and the inter-ring Region 2 (right panel). The magnitudes from bottom to top are $BVIUR$. } \\label{fig03} \\end{figure}  Significant dust formed in SN1987A starting around day 500, and by day 1000, the bolometric flux of the nebula was dominated by dust radiation with a temperature of $T_{dust}\\sim200$K and mid-infrared line emission(\\cite{Sun_etal92,Woo_etal93}). In Fig.~\\ref{fig04} I show the latest data on the mid-infrared evolution of SN1987A with the ISOCAM detection (\\cite{Fis_etal02}) and a detection from OSCIR on the CTIO 4m telescope (\\cite{Bou02}) at 10\\umu{m}. Without 20\\umu{m} detections however, we can't estimate the dust temperature and cannot measure the dust emission. If we assume the same bolometric corrections from day 1800, we find the bolometric flux of SN1987A is ${\\rm{log}_{10}(L)\\sim36.0}$ erg s$^{-1}$.  \\begin{figure}[h] \\begin{center} \\includegraphics[angle=0,width=0.7\\textwidth]{nsuntzeff_fig04.eps} \\end{center} \\caption[nsuntzeff_fig04.eps]{Mid-infrared light curves of SN1987A} \\label{fig04} \\end{figure} ",
        "conclusions": ""
    },
    "0212/astro-ph0212427_arXiv.txt": {
        "abstract": "We present a detailed comparison of structural properties in the rest-frame $V$-band of cluster and field galaxies, selected and analyzed in the same manner, to test the hypothesis that much of the current cluster galaxy membership resulted from the fairly rapid (1-2 gigayears) transformation of infalling, field spirals into red, cluster early-types.  Specifically, we have selected $\\sim140$ galaxies from three nearby Abell clusters (A85, A496 and A754) that have colors significantly bluer than the red sequence population, and compared them to $\\sim80$ field galaxies with similar colors and luminosities from Jansen et al. (2000, ApJS, 126, 271). The comparison is based on the hypothesis that recent (1-4 gigayears) cluster arrivals were originally blue and star-forming, then stopped forming stars to dim and redden in a few gigayears. For the comparison we quantify galaxy internal structure and morphology from two-dimensional bulge/disk decompositions using GIM2D.  We observe structural differences between blue galaxies in local ($z<0.06$) clusters, compared to field environments.  All cluster galaxies have spectroscopic membership. The majority of blue cluster members, presumably recent additions, are physically smaller and fainter than their equally-colored field counterparts.  At a matched size and luminosity, the newer cluster arrivals are quantifiably smoother in appearance, yet their total light is as disk-dominated as in normal field spirals. Moreover, half of the blue cluster members appear to have blue cores or globally blue color profiles, in contrast with field spirals, which typically exhibit red inward color gradients. Blue cores suggest enhanced nuclear star formation, possibly a starburst, while uniformly blue profiles are consistent with an episode of fairly strong global star formation in the past few gigayears.  Our previous work (McIntosh et al. 2004, ApJ submitted) shows that the blue membership of local clusters is a recently infalling population that has yet to encounter the dense core. In a Universe {\\it without environmental dependent evolution outside of cluster cores}, we would expect blue disk galaxies inhabiting field and cluster regions to have similar morphology, size, and color gradient distributions. Our findings show conclusively that not only the abundance of red and blue galaxies depends on environment, but also that fundamental structural and morphological galaxy properties do indeed reflect the environment in which the galaxy is found. Moreover, the data show that the transformation of accreted galaxies is {\\it not} confined to the dense cluster core. The overall properties of bluer cluster members are best explained by environment-driven transformation of accreted field spirals, and our results suggest that the processes that govern color and morphological evolution occur separately. ",
        "introduction": "If structure in the Universe forms hierarchically, the massive galaxy clusters we observe locally should have been built through the infall and accretion of matter ({\\it i.e.} galaxies, groups, subclusters) along large-scale filaments \\citep[][and references therein]{white78,white91,west95,kodama01c} As new galaxies fall into regions of greater density, they must interact with both the hot, inter-cluster medium (ICM) and the increased numbers of nearby, cluster neighbors. Given these interactions, it seems plausible that infalling populations undergoing ``normal'' star formation rates (SFR) will be transformed into red lenticular (S0) galaxies \\citep[hereafter the ``transformation scenario'',][]{gunn72,charlot94,abraham96,dressler97,vandokkum98,moore99,poggianti99}. The transformation scenario has been invoked to explain the well-known evolution of cluster galaxy color \\citep[the Butcher-Oemler effect, \\eg][]{bo78a,bo84,rakos95,margoniner01,ellingson01} and morphology \\citep{lavery88,couch94,dressler94a,dressler94b,ellis97,oemler97}. This scenario is key to the idea of ``progenitor bias'' put forth by \\citet{vandokkum01}, which states that the progenitors of the youngest early-type galaxies in clusters were morphologically transformed from later types, and thus would not be considered as part of the early-type population of clusters at larger redshifts. Moreover, this environmental based scenario is compatible with the observed correlations between increased density and cluster galaxy morphology \\citep[the ${\\rm T}-\\Sigma$ relation, \\eg][]{melnick77,bo78b,dressler80,postman84,dressler97,hashimoto99} and star formation (SF) properties \\citep{balogh97,balogh98,hashimoto98,lewis02}.  Establishing whether the transformation scenario is occurring in low redshift clusters provides important confirmation for the continued hierarchical evolution of these massive systems to the present epoch.  For example, finding cluster members in the midst of this predicted evolution will secure the present-day whereabouts of the blue galaxies that were more abundant in $z\\sim0.3$ clusters (Butcher-Oemler effect), and may give us additional clues to how and where galaxies evolve in dense environments. The literature is replete with evidence for the evolution and environmental dependence of cluster galaxy properties.  Yet, even with the huge body of observational data on galaxy clusters in the local Universe, there has been no direct confirmation showing any fraction of their members in the midst of this predicted galaxy transformation.  Therefore, this lack of confirmation means either that such transformations are no longer ongoing in the present-day Universe, or that the data and observational tests have been insufficient.  To this end we have wide-field imaged three nearby Abell clusters with hundreds of spectroscopic member galaxies to look for evidence of the transformation scenario.  In \\citet[][hereafter Paper~2]{mcintosh04} we established the existence of significant fractions ($18-23\\%$) of blue and moderately blue members with kinematic and spatial properties as expected for a recently infalling population \\citep{diaferio01}.  Finding these populations of late arrivals provided one crucial test of the validity of the transformation scenario.  Moreover, we found that many of the blue members have disky visual morphologies, yet lack strong spiral features, perhaps the result of morphological smoothing. In this paper we use quantitative measures of morphology and structure to determine whether the appearances of the recent cluster arrivals differ from those of a control sample of star-forming spirals from the lower density ``field'' environment. Physical differences between cluster members and field galaxies with similar star-forming blue colors will provide direct observational confirmation for the morphological smoothing aspect of the transformation scenario.  Various physical mechanisms have been put forth to explain the transformation scenario in clusters. There is general agreement that the denser environment will have an effect on the SF properties of new arrivals.  The cold neutral gas supply may be removed by ram-pressure stripping during a galaxy's virial motion through the hot ICM \\citep{gunn72,solanes92,abadi99,quilis00}, or tidal stripping by either the cluster potential or individual galaxies \\citep{spitzer51,valluri90}. Alternatively, the star-forming gas could be rapidly consumed in a burst of SF triggered by ram-pressure effects \\citep{dressler83,fujita99}, tidal effects \\citep{byrd90,hashimoto98,bekki99}, or galaxy merging \\citep{barnes91,bekki98}. Other authors have suggested that the replenishment of star-forming gas is cut off instead, and that the continuing formation of stars gradually ($>1$~gigayear) exhausts the remaining supply of \\ion{H}{1} \\citep{larson80,kodama01a,balogh00,balogh00b}. Whether the SF fuel supply is removed, used up or cut off, the result is the eventual end of SF and the passive fading and reddening of the once blue galaxies as their hot, young stars die and their mean stellar population age grows older.  This passive ``color evolution'' has a minimum timescale equal to the lifetimes of the short-lived luminous stars ({\\it i.e.} 1-2~gigayears for A-type stars), and longer if SF is slowly diminished, rather than suddenly truncated.  It stands to reason that the cessation of SF in spirals will eventually produce an overall smoother galaxy profile by reducing substructure ({\\it e.g.} spiral arms, \\ion{H}{2} knots), yet this may take several gigayears or longer.  Some dynamical simulations predict that galaxies falling into dense clusters will undergo more rapid ($\\sim1$~gigayear) changes in their physical morphology due to encounters with other cluster members.  Among processes that predict rapid spiral to S0 evolution are ``galaxy harassment'', the smoothing of galaxy appearances and stripping of matter from multiple tidal encounters with other cluster members \\citep{moore96,moore96b,moore98,moore99}, and unequal-mass spiral mergers \\citep{bekki01,cretton01}. Other processes, such as ``passive'' spiral formation from halo gas stripping and subsequent ``starvation'', predict a much more gradual transition ($\\sim3$ gigayears) into S0s \\citep{bekki02}. Presently, it is unclear whether color and physical morphology evolution occur simultaneously, or if they are decoupled \\citep{poggianti99,couch01,kodama01b}.  Typical Abell clusters have measured velocity dispersions of $\\sigma\\leq1000$~\\kms\\ \\citep[\\eg][]{depropris02}, which correspond to roughly 1~Mpc per gigayear. At this rate, recent arrivals would take several gigayears to make their way from the cluster outskirts to the inner regions and, given the timescales for the various proposed mechanisms, much of the predicted rapid evolution would have already transpired. Furthermore, the extreme environments within the inner half megaparsec of clusters will destroy or seriously truncate disk galaxies \\citep{moore98}, hence looking for recent arrivals and evidence of transformation towards the centers of clusters is not profitable. Unfortunately, until quite recently CCD imagers were limited to small fields of view so that it was quite difficult to image large angular regions of the sky.  Thus, the bulk of cluster imaging work in the nearby Universe ($cz<15,000$~\\kms) has concentrated on galaxies belonging to the core.  For example, at the distance to the well-studied Coma cluster, 2~Mpc in diameter is projected across two degrees of arc on the sky.  \\citet{terlevich01} used overlapping CCD frames to study a square-degree region of Coma that still only provided a view of the inner $R=0.5$~Mpc.  They found no evidence for a population of recent arrivals in mid-transformation.  Yet, observations of more distant ($z>0.3$) clusters have yielded cluster S0s with blue colors -- possible mid-evolution examples of transforming blue spirals into red S0s \\citep{vandokkum98,rakos00,smail01}.  Therefore, our observing campaign has taken advantage of recent wide-field imaging capabilities and large redshift surveys to explore galaxies at outer ($R>0.5$ Mpc) regions in three local Abell clusters: A85 ($z=0.055$), A496 ($z=0.033$) and A754 ($z=0.055$). Most cluster studies rely upon statistical corrections to remove foreground and background galaxy contaminates; therefore, cluster membership using such corrections becomes more uncertain at larger cluster radii.  Using the large cluster galaxy redshift survey of \\citet{christlein03}, we have limited our study to only galaxies confirmed as members with spectroscopic redshifts within $\\pm3\\sigma$ of the mean cluster redshift.  To test the transformation scenario we look for structural differences between blue galaxies residing in cluster and field environments through a detailed analysis of their $V$-band structural properties from two-dimensional bulge/disk (B/D) decompositions using GIM2D \\citep{simard02}. We assume that infalling galaxies suffer SF truncation and evolve rapidly in color and luminosity; hence, relative blueness gives us an estimate of cluster accretion time and ``membership age''.  In Paper 2 we performed a detailed $U,V$ color-magnitude relation (CMR) analysis for 637 galaxies belonging to the three clusters in our sample.  The bluer members reside preferentially in the cluster outskirts (typically $R>0.5$~\\hmpc) and have kinematics suggesting a non-virialized, infalling population \\citep{diaferio01}. For this study we select the 143 members with $M_V\\leq -17+5\\log{h}$ and colors more than $2\\sigma_{\\rm CMR}$ bluer than their cluster CMR, where $\\sigma_{\\rm CMR}$ is the measured CMR scatter ($<0.1$~mag) in $(U-V)$ color (Paper~2).  For comparison, we select 78 galaxies with similar colors and luminosities from the Nearby Field Galaxy Survey (NFGS) of \\citet{jansen00}.  We degrade and resample the field members to measure their properties as if they were observed at $z=0.055$.  In the transformation scenario we expect that the most recent cluster arrivals will have bluish colors from recent (within $<2$ gigayears) SF, disk-like morphologies reflective of an infalling spiral population, and smoother appearances sometime after SF cessation.  In Paper 2 we find most blue members in these clusters are classified visually as disk-dominated systems with weak spiral features. SF and morphology evolution may be decoupled in cluster galaxies \\citep[e.g.][]{couch01}, and whether or not the newest (blue) members already exhibit morphological smoothing provides a key motivation for this work. Furthermore, processes like galaxy harassment will reduce the relative disk substructure and size \\citep{moore98}, and a variety of mechanisms predict centralized starbursts. In our third paper (hereafter Paper~3, D. H. McIntosh, H.-W. Rix, \\& N. Caldwell, in preparation) we establish $U$ and $V$-band structural differences between three color-magnitude (C-M) selected cluster populations.  The bluer membership found typically at large cluster-centric distances consists of predominantly disky systems with uniformly blue to blue centered profiles, lending additional credence to their recent infall origin.  Of interest, we notice that the blue members have little excess substructure relative to the redder, presumably older members with smooth morphologies.  By looking for structural differences between cluster and field blue galaxies, we can determine whether recent cluster arrivals have undergone transformation compared to their presumed progenitors, star-forming galaxies believed to be the infall source in an hierarchical cosmology.  Here we concentrate our effort on $V$-band properties given the higher signal-to-noise (S/N) and better seeing characteristics of our cluster imaging in this passband.  In this paper we present evidence for environment-driven galaxy transformation through a detailed comparison of cluster members and field galaxies with similar luminosities and blue colors.  We briefly summarize the cluster and field blue galaxy sample selections in \\S2. In \\S3 we describe our two-dimensional B/D decompositions to measure quantitative morphologies in both the cluster and field samples. We include a detailed study of the limitations of this method to yield useful morphological measurements from the cluster and field data. In \\S4 we analyze detailed comparisons for the observed properties of cluster members against field counterparts selected in the same manner.  We discuss our results in light of morphological transformation processes in \\S5. And we give our conclusions in \\S6. Throughout this paper we use $h = H_0/(100$~km~s$^{-1}$~Mpc$^{-1})$, and we assume a $\\Lambda$-CDM (cold dark matter), $\\Omega_{\\rm M}=0.3$ and $\\Omega_{\\rm \\Lambda}=0.7$, flat ($\\Omega_{\\rm k}=0$) cosmology. ",
        "conclusions": "The morphological transformation of field (or filament) spirals during infall into the dense cluster environment \\citep[\\eg][]{kodama01b} has been proposed to explain the rapid evolution of cluster galaxies over the past 5~gigayears. In particular, the spiral-to-S0 ratio in clusters has decreased from $\\sim2:1$ at $z=0.5$ to $\\sim1:2$ at present day, suggesting a direct evolutionary link between the decreasing numbers of spirals and increasing numbers of S0s \\citep{dressler97}.  In this paper we uncover evidence for environment-driven galaxy transformation through a detailed comparison of recent cluster arrivals with galaxies of similar luminosities and blue colors in the field.  In a Universe {\\it without environmental dependent evolution outside of the dense cluster cores}, we would expect blue disk galaxies inhabiting field and cluster regions to have similar morphology, size, and color gradient distributions. Our findings show conclusively that fundamental galaxy properties do indeed reflect the environment in which the galaxy is found.  We find structural differences between the blue galaxies inhabiting nearby ($z<0.06$) clusters, compared to field environments.  The majority of blue cluster members are physically smaller and fainter than their field counterparts.  At a matched size and luminosity, the newer cluster arrivals have quantifiably less internal substructure, yet have equally disk-dominated morphologies as normal field spirals. Furthermore, half of blue cluster galaxies have blue cores or globally blue color profiles in contrast with field spirals which typically show redward color gradients. Blue cores suggest enhanced nuclear star formation, possibly a starburst, while uniformly blue profiles are consistent with an episode of fairly strong global star formation in the past few gigayears.  We show in Paper~2 that blue cluster galaxies are members of a recently infalling population that has not yet encountered the cluster core.  Therefore, the differences we observe between very blue cluster and field galaxies show that galaxy transformation occurs in accreted galaxies {\\it before} the violent effects (e.g. strong tides, ram-pressure) of the cluster core come into play.  Much of what we observe in the blue galaxy populations of local clusters can be explained by the process of galaxy harassment \\citep{moore96,moore96b,moore98,moore99}, under several imposed conditions. Nevertheless, ram-pressure stripping \\citep{gunn72,abadi99,quilis00} must play an important role in assuring that SF is eventually quenched in cluster galaxies to produce the strong color evolution, especially towards the cluster core, implied by the Butcher-Oemler effect \\citep[][and references therein]{bo84,rakos95,margoniner01}. The fact that transformation has occurred in cluster galaxies with spiral-like colors (the VBGs) shows that the processes that govern color (SF) and physical morphology evolution are decoupled."
    },
    "0212/astro-ph0212341_arXiv.txt": {
        "abstract": "HII regions are known to contribute to the so-called {\\em thin layer} of the diffuse Warm Ionized Gas. In order to constrain this contribution, we reconstruct the 3-D distribution of the sources. A detailed spatial analysis of the largest up-to-date sample of HII regions is presented. ",
        "introduction": "The thin disk component of the Galactic free electron distribution is largely contributed by localized HII regions. In order to constrain this contribution, it is important to reconstruct the spatial distribution of known sources.\\\\ In a recent paper (Paladini et al., 2002, hereafter Paper I), we describe the construction of an extensive radio catalog (1442 sources) of Galactic HII regions by the combination of 24 published lists and catalogs of these objects. The final compilation consists of a Master Catalog (containing original data and corresponding errors) and a Synthetic Catalog at 2.7 GHz (which summarizes the basic information - flux density, angular diameter and $V_{LSR}$ - for each source of the Master Catalog). The Synthetic Catalog has provided the source of information for the spatial analysis here presented. ",
        "conclusions": ""
    },
    "0212/astro-ph0212388_arXiv.txt": {
        "abstract": "s{ A fit to  the present cosmic ray positron fraction can be considerably improved, if in addition to the positron production by  nuclear interactions in the universe the possible contribution from supersymmetric dark matter annihilation is taken into account. We scan over the complete SUSY parameter space of the Constrained Minimal Supersymmetric Model (CMSSM) and find that in the acceptable regions the neutralino annihilation into $b\\overline{b}$ quark pairs is the dominant channel with hard positrons emerging from the semileptonic decays of the B-mesons. } ",
        "introduction": "The cosmic ray positron fraction at momenta above 7 GeV is difficult to describe by the background only hypothesis\\cite{HEAT}. A contribution from the annihilation of neutralinos, which are the leading candidates to explain the cold dark matter in the universe, can improve the fits considerably\\cite{edsjo,kane,deboerastro}. The neutralinos are the Lightest Supersymmetric Particles (LSP) in  supersymmetric extensions of the Standard Model, which are stable, if R-parity is conserved.  This new multiplicative quantum number for the supersymmetric partners of the Standard Model (SM) particles is needed to prevent proton decay and simultaneously prevents the LSP a) to decay into the lighter SM particles and b) can only interact with normal matter by producing additional supersymmetric particles. The cross sections for the latter are typically of the order of the weak cross sections, so the LSP is ``neutrinolike'', i.e. it would form halos around the galaxies and consequently, it is an excellent candidate for dark matter.  In addition to being of interest for cosmo\\-logy, supersymmetry solves also many outstanding problems in particle physics, among them\\cite{rev}: 1) it provides a  unification of the strong and electroweak forces, thus being a prototype theory for a Grand Unified Theory (GUT) 2) it predicts spontaneous electroweak  symmetry breaking (EWSB) by radiative corrections through the heavy top quark, thus providing a relation between the GUT scale, the electroweak scale and the top mass, which is perfectly fullfilled 3) it includes gravity 4) it cancels the quadratic divergencies in the Higgs mass in the SM 5) the lightest Higgs mass can be calculated to be below 125 GeV in perfect agreement with electroweak precision data, which prefer indeed a light Higgs mass. Therefore it is interesting to study the positron fraction from neutralino annihilation in the reduced region of SUSY parameter space, where these constraints are satisfied and compare it with data, as  will be done in the next section. It should be noted that the positrons from  the annihilation of heavy particles leads to a much harder energy spectrum than the positrons expected from nuclear interactions, thus leading to a significant change in the shape of the spectrum. In addition, in the positron fraction, defined as $e^+/(e^- +e^+)$, many systematic uncertainties related to fluctuations in the nuclear densities and propagation in the universe, cancel. \\begin{figure} \\begin{center} \\includegraphics [width=0.49\\textwidth,clip]{mass-evol.htb.eps} \\includegraphics [width=0.49\\textwidth,clip]{gfrac50_sf.eps} \\caption[]{\\label{mass} \\it Left: The running of the masses between the GUT scale and the electroweak scale for \\tb=35. Right: The gaugino fraction as function of $m_0$ and $m_{1/2}$ for $\\tan\\beta=35$. For lower va\\-lues of \\tb the gaugino fraction is even closer to one. The region with gaugino fraction zero is excluded, since the LSP would be charged, which is excluded because the dark matter is neutral. } \\end{center} \\end{figure} ",
        "conclusions": ""
    },
    "0212/astro-ph0212177_arXiv.txt": {
        "abstract": "We estimated the rate of comet and asteroid collisions with the terrestrial planets by calculating the orbits of 13000 Jupiter-crossing objects (JCOs) and 1300 resonant asteroids and computing the probabilities of collisions based on random-phase approximations and the orbital elements sampled with a 500 yr step. The Bulirsh-Stoer and a symplectic orbit integrator gave similar results for orbital evolution, but may give different collision probabilities with the Sun. A small fraction of former JCOs reached orbits with aphelia inside Jupiter's orbit, and some reached  Apollo orbits with semi-major axes less than 2 AU, Aten orbits, and inner-Earth orbits (with aphelia less than 0.983 AU) and remained there for millions of years. Though less than 0.1\\% of the total, these objects were responsible for most of the collision probability of former JCOs with Earth and Venus. We conclude that a significant fraction of near-Earth objects could be extinct comets that came from the trans-Neptunian region. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/astro-ph0212031_arXiv.txt": {
        "abstract": "This paper discusses three new issues that necessarily arise in realistic attempts to apply nonlinear dynamics to galaxy evolution, namely: (i) the meaning of chaos in many-body systems, (ii) the time-dependence of the bulk potential, which can trigger intervals of {\\em transient chaos}, and (iii) the self-consistent nature of any bulk chaos, which is generated by the bodies themselves, rather than imposed externally. Simulations and theory both suggest strongly that the physical processes associated with galactic evolution should also act in nonneutral plasmas and charged particle beams. This in turn suggests the possibility of testing this physics in real laboratory experiments, an undertaking currently underway. ",
        "introduction": "As recently as 1990, most galactic dynamicists ignored completely the possible role of chaos in galaxies. However, the past decade has seen a growing recognition that chaos can be important in determining the structure of real galaxies. Still, much recent work involving chaos in galactic astronomy has been simplistic in that it has involved naive applications of standard techniques from nonlinear dynamics developed to analyse two- and three-degree-of-freedom time-independent Hamiltonian systems. The object here is to discuss some of the additional complications which arise if nonlinear dynamics is to be applied to real galaxies, which are {\\em many-body systems} comprised of a large number of interacting masses and characterised by a {\\em self-consistently determined bulk potential} which, during their most interesting phases, can be {\\em strongly time-dependent.} ",
        "conclusions": "This paper has focused on several fundamental issues that arise in attempts to apply nonlinear dynamics to real galaxies, many-body systems characterised by a self-consistently determined bulk potential which, during their most interesting phases, can be strongly time-dependent. As recently as a decade ago these issues would have been considered of largely academic, rather than practical, interest. However, recent observational advances -- which facilitate improved high resolution photometry of individual objects as well as statistical analyses of large samples with varying redshift -- and improved computational resources -- which allow unparalleled explorations of multi-scale structure --, together with the recognition that the basic physics can be also probed in the context of charged particle beams, imply that theoretical predictions regarding these `academic' issues can in fact be tested observationally, numerically, and experimentally. \\vskip .1in \\par\\noindent Thanks to Ioannis Sideris for providing Fig.~1 and to Rami Kishek for providing Fig.~3. Thanks also to Court Bohn for his comments on a preliminary draft of the manuscript. Limited financial support was provided by NSF AST-0070809."
    },
    "0212/astro-ph0212540_arXiv.txt": {
        "abstract": "We investigate the effect that the structure of GRB jets has on the afterglow light curves for observers located at different viewing angles, $\\theta_{\\rm obs}$, from the jet symmetry axis. The largest uncertainty in the jet dynamics is the degree of lateral energy transfer. Thus, we use two simple models, that make opposite and extreme assumptions for this point, and calculate the light curves for an external density that is either homogeneous, or decreases as the square of the distance from the source. The Lorentz factor and kinetic energy per unit solid angle are initially taken to be power laws of the angle from the jet axis. We perform a qualitative comparison between the resulting light curves and afterglow observations. This constrains the jet structure, and poses problems for a `universal' jet model, where all GRB jets are assumed to be intrinsically identical, and differ only by our viewing angle, $\\theta_{\\rm obs}$. ",
        "introduction": "\\label{sec:intro}  There are several different lines of evidence in favor of collimated outflows, also referred to as jets, in gamma-ray bursts (GRBs). Perhaps, the best evidence so far is the achromatic break seen in the light curves of many (though not all) GRB afterglows. A collimated outflow also helps to reduce the estimate for the total energy output in gamma-rays that is inferred from the fluence in GRBs with known redshifts, which in some cases approaches and in one case (GRB 990123) even exceeds the rest energy of a solar mass star, {for a spherically symmetric emission}. Such an energy output in gamma-rays is hard to produce in any model that involves a stellar mass object. Furthermore, a non-spherical flow can also manifest itself through linear polarization (Sari 1999; Ghisellini \\& Lazzati 1999) as was indeed observed for a few afterglows (Covino et al. 1999; Wijers et al. 1999; Rol et al. 2000).  Most GRB jet models consider an outflow that is uniform within some finite, well defined, opening angle around its symmetry axis, and where the Lorentz factor, energy density etc. drop sharply beyond this opening angle (Rhoads 1997, 1999; Panaitescu \\& M\\'esz\\'aros 1999; Sari, Piran \\& Halpern 1999; Kumar \\& Panaitescu 2000; Moderski, Sikora \\& Bulik 2000; Granot et al. 2001, 2002). Such a uniform jet with a sharp, well defined, edge shall be referred to as a `top hat' jet. The possibility that GRB jets can display an angular structure, where the kinetic energy per unit solid angle, $\\epsilon$, and the Lorentz factor, $\\Gamma$, of the GRB outflow vary smoothly as power laws in the angle, $\\theta$, from the jet axis, was proposed by M\\'esz\\'aros, Rees \\& Wijers (1998). We shall refer to such smoothly varying relativistic outflows as `structured' jets (as opposed to a `top hat' jet, that has no inner structure).  Recently, several different groups have analyzed afterglow observations within the frame work of the `top hat' jet model, and have inferred a relatively narrow distribution both for the total energy output in gamma-rays (Frail et al. 2001) and in the initial kinetic energy of the relativistic outflow (Panaitescu \\& Kumar 2001; Piran et al. 2001). These results may alternatively be interpreted as GRB jets having a universal structure, which is intrinsically the same for all GRBs, and the observed differences between different GRBs are a result of different viewing angle, $\\theta_{\\rm obs}$, w.r.t. the jet symmetry axis (Lipunov, Postnov \\& Prokhorov 2001; Rossi, Lazzati \\& Rees 2002; Zhang \\& M\\'esz\\'aros 2002). Whereas in the `top hat' jet interpretation, the jet break time, $t_j$, depends mainly on the initial opening angle of the jet, $\\theta_j$ (and also has a smaller dependence on its energy per unit solid angle and on the external density, e.g. Sari, Piran \\& Halpern 1999), in the universal `structured' jet interpretation, $t_j$ depends mainly on the viewing angle, $\\theta_{\\rm obs}$, and the light curve is similar to that for a `top hat' jet with an opening angle $\\theta_j=\\theta_{\\rm obs}$  and the same value of the energy per unit solid angle of the `structured' jet at $\\theta=\\theta_{\\rm obs}$.  While the evolution of `top hat' jets and their light curves have been widely investigated (Rhoads 1999; Panaitescu \\& M\\'esz\\'aros 1999; Kumar \\& Panaitescu 2000; Moderski, Sikora \\& Bulik 2000; Granot et al. 2002), including numerical simulations of the jet dynamics and calculation of the resulting light curves (Granot et al. 2001), much less work has been done on `structured' jets. In an accompanying paper (Kumar \\& Granot 2002) we calculate the dynamics of `structured' relativistic jets by solving relativistic fluid dynamics equations, and demonstrate that a simple analytic or semi-analytic model can qualitatively reproduce the results for the jet dynamics and for the light curves as long as the Lorentz factor along the jet axis is of order 4 or larger. In this paper we use two simple models for the jet dynamics, which are designed to bracket the true hydrodynamic evolution of the jet, to calculate the afterglow light curves for a wide range of parameters for the structured jet and the external density. A qualitative comparison of the light curves with afterglow observations provides constraints on the jet structure and density profile in its immediate vicinity.  In \\S \\ref{sec:model} we present the physical model. We begin with the initial conditions (\\S \\ref{IC}), then describe our two simple dynamical models (\\S \\ref{JD}) and finally outline the procedure for calculating the light curves (\\S \\ref{LC}). Our results are presented in \\S \\ref{sec:results}, and in \\S \\ref{sec:dis} we discuss our main conclusions. ",
        "conclusions": "\\label{sec:dis}  We have calculated light curves from structured relativistic jets (jets whose Lorentz factor and energy per unit solid angle vary smoothly with angle from the jet axis) using two simple models for the jet dynamics, motivated by the hydrodynamical simulation of Kumar and Granot (2002), and as two limiting cases of energy redistribution in the lateral direction.  The first model considers the energy per unit solid angle, $\\epsilon(\\theta,t)$, to be time independent, so that each segment of the jet evolves independently of the other parts of the jet, as if it was part of a spherical flow. This is roughly consistent with the results from hydrodynamical simulations of structured jets for which $\\epsilon$ varies slowly with $\\theta$, where we find the transverse velocity in the comoving frame to be small compared to the sound speed throughout much of the jet evolution, as long as the Lorentz factor, $\\Gamma$, is of order a few or larger along the jet axis (Kumar \\& Granot 2002). This model is expected to be a good approximation for calculating the jet dynamics and the resulting afterglow light curves for the first few days after a GRB.  The second model considers the maximum possible re-distribution of energy in the transverse direction so as to reduce the lateral-gradient of $\\epsilon$. The energy per unit solid angle, $\\epsilon(\\theta,t)$, is taken to be proportional to the average over the initial distribution of $\\epsilon$, $\\epsilon(\\theta,t_0)$, within the area out to which a sound wave could have traveled from $\\theta$ since the initial time, $t_0$. This model is likely to be more accurate in describing the late time behavior of the jet and the light curves, when the transverse velocity becomes of order the sound speed.  We have calculated light-curves for a number of different initial jet structures, which are taken to be power-law profiles at angles larger than some core angle, i.e. $\\epsilon\\propto \\theta^{-a}$ and $\\Gamma\\propto \\theta^{-b}$ for $\\theta>\\theta_c$. We have considered three different jet profiles, namely $(a,b)$ = (2,0), (2,2), (0,2) for both a homogeneous external medium as well an ambient density falling off as distance squared from the center of explosion (i.e. $k=0$ or $2$, where $\\rho_{ext}\\propto r^{-k}$).  The quantitative differences in the light-curves calculated using our two models are typically of order unity for all of the jet structures we have considered. This gives us some confidence that lightcurves can be calculated with a reasonable accuracy even with a crude modeling of jet dynamics. There are, however, interesting qualitative differences in the light-curve properties depending on how we model jet dynamics. For instance, for $(a,b)=(2,2)$ and a homogeneous medium ($k=0$), the temporal index of the light curve, $\\alpha$ (defined by $F_\\nu\\propto t^\\alpha$), before the jet break is larger for large viewing angles, $\\theta_{\\rm obs}$. This effect is much more prominent for model 1, compared to model 2. For $(a,b)=(2,0)$ and $k=0$, with model 1 there is a pronounced flattening in the light curve just before the jet break for $\\theta_{\\rm obs}\\gtrsim 3\\,\\theta_c$, while for model 2 this effect is much smaller.  There is a significant difference between the two models in the late time light curves, as we expect: model 1 light curves continue a power-law decline at late times, whereas model 2 light curves show a slight flattening at late times resulting from energy redistribution -- see Figures \\ref{LCk0} and \\ref{alpha_k0} for homogeneous external medium with jet structures $(a,b)=(2,2)$ \\& (2,0), and Figures \\ref{LCk2} and \\ref{alpha_k2} for the stratified external medium the particular case of $(a,b)=(2,0)$. Figures \\ref{LCk2} and \\ref{alpha_k2} also show that the break in the light curves for a stratified external medium is very gentle and the change to the power-law index $\\alpha$ is small so long as the observer is not too far the axis: $\\theta_{\\rm obs}/\\theta_c\\lesssim 10$. This result is consistent with the work of Kumar \\& Panaitescu (2000), based on the analysis of a `top hat' jet model, that it is very difficult to see a break in light curves for jets in an $r^{-2}$ medium.  We can learn about the jet structure in GRBs from comparing our theoretical light curves to GRB afterglow observations. For instance, the light curves for $(a,b)=(2,0)$ and (2,2) behave very differently prior to the jet break. This can be used to discriminate between these possibilities (Figures \\ref{LCk0} \\& \\ref{alpha_k0}). For obvious reasons, the light curves for the (0,2) case are very distinct and similar to a spherically symmetric explosion, and do not show a jet break. This jet structure can therefore be ruled out for all GRBs that show a jet break in their afterglow light curves. Clear breaks in the light curves can be seen for a stratified external medium density only when the observer is located far away from the jet axis ($\\theta_{\\rm obs}\\gtrsim 10\\,\\theta_c$), however, in this case the GRB will be faint and is likely to be missed in flux limited triggers for detecting GRBs.  The light curves we obtain for structured jets, in many cases show a different qualitative behavior compared to afterglow observations. For example, $k=2$ does not produce a sufficiently sharp jet break, while for $k=0$ only $a\\approx 2$ produces jet breaks that resemble afterglow observations, and even then, it is very difficult to produce all current afterglow observations with a universal jet profile, that is, if all GRB jets have the same structure and differ only by our viewing angle, $\\theta_{\\rm obs}$, w.r.t. the jet symmetry axis. Even if GRB jets do not have a universal structure, a comparison between the observed light curve and theoretical calculations can constrain the jet structure of each GRB separately.  Future afterglow observations should enable us to determine the structure of relativistic jets in GRBs as well as the properties of the surrounding medium. This will enable us to determine if the energy release in GRBs is nearly constant or not, and whether the observed afterglow light curves and the jet break time are determined by the observer angle w.r.t. a structured jet, or by the opening angle of a `top hat' jet."
    },
    "0212/astro-ph0212156_arXiv.txt": {
        "abstract": "Though widely observed to be emanating from a variety of astrophysical sources, the underlying physical mechanism behind the formation of galactic and extragalactic outflows is still enshrouded in a veil of mystery.  In addition, it has not been possible to calculate accurately the amount of matter expelled in these events.  In this article we present a {\\em non-self-similar analytical model}, which, for the first time, we believe is able to explain the outflow formation phenomenon as well as compute the mass outflow rate by simultaneously solving the equations governing the exact transonic accretion and outflow.  Our model predicts the dependence of this rate on various flow parameters as well as the {\\em exact} location from where the outflows are launched. ",
        "introduction": "\\subsection{Model Description and Prescription for the Formation of Outflow Generating Surface} We consider thin, axisymmetric polytropic inflows around a Scwarzschild Black hole in vertical equilibrium.  We ignore the self-gravity of the flow and calculations are done using Paczy\\' nski-Wiita (Paczyn\\'ski \\& Wiita, 1980) potential which mimics surroundings of the Schwarzschild black hole. Considering the inflow to be polytropic, we explore both the polytropic and the isothermal outflow.\\\\ Due to the fact that close to the BH the radial component of the infall velocity of accreting material would be enormously high, viscous time scale would be much longer than the infall time scale and a rotating inflow entering into a black hole  will have almost constant specific angular momentum close to the black hole for any moderate viscous stress (Das, 1998 and DC99). Though at the outer edge of the accretion disk the angular momentum distribution may be Keplerian or even super-Keplerian, matter would be highly sub-Keplerian close to the black hole to satisfy the inner boundary condition at the event horizon of the BH (see Chakrabarti, This volume and references therein). This almost constant angular momentum produces a very strong centrifugal force  which increases much faster compared to the gravitational force and becomes comparable at some specific radial distance, location of which is easy to compute. Here, (actually, a little farther out, due to thermal pressure) matter starts piling up and produces the centrifugal pressure supported boundary layer (CENBOL). Further close to the black hole, the gravity always wins and matter enters the horizon supersonically after passing through a sonic point. Formation of CENBOL may be attributed to the shock formation in accreting fluid or to the maximisation of polytropic pressure of the inflow. In CENBOL region the flow becomes hotter and denser and for all practical purposes behaves as the stellar atmosphere so far as the formation of outflows are concerned. A part of the hot and dense shock-compressed inflowing material is then `squirt' as outflow from the CENBOL. In case where the shock does not form, regions around pressure maximum achieved just outside the inner sonic point of the {\\it inflow} would also drive the flow outwards. The outflow is shown to be thermally and centrifugally accelerated but is assumed to be confined by external pressure of the ambient medium.  Subsonic outflows originating from CENBOL would pass through sonic points and reach far distances as in wind solution.\\\\ It is interesting to `visualize' how the combined accretion-outflow system along with the central accretor would `look like' in reality. \\begin{figure} \\vbox{ \\vskip -0.0cm \\centerline{ \\psfig{figure=jetdisc.ps}}} \\caption[]{Multicomponent combined flow geometry in 3-Dimension.} \\label{penG} \\end{figure} In the following figure we attempt to illustrate the 3-D geometry of coupled disk-outflow system according to our model. {\\bf B} is the accreting Schwarzschild black hole while {\\bf C}'s represent the hot and dense CENBOL region. {\\bf D(K)} and {\\bf D(SK)} represent the thin Keplarian part and puffed-up sub-Keplarian part of the advective accretion disk respectively (measurements not in scale). Due to the axisymmetry assumption in accretion, two oppositely directed jet are expelled from the close viscinity of {\\bf B}. The inner and the outer surfaces of the outflow are the funnel wall and centrifugal barrier respectively as explained in Das, 1998 and in DC99. \\subsection{The Overall Solution Scheme} Let us suppose that matter first enters through the outer sonic point and passes through a shock. At the shock, part of the incoming matter, having higher entropy density is likely to return back as winds through a sonic point, other than the one it just entered. Thus a combination of topologies, one from the region of accretion and the other from the wind region is required to obtain a full solution. In the absence of the shocks, the flow is likely to bounce back at the pressure maximum of the inflow and since the outflow would be heated by photons, and thus have a smaller polytropic constant, the flow would leave the system through an outer sonic point different from that of the incoming solution. By simultaneously solving the proper set of equations in appropriate geometry (see Das, 1998, DC99 and Chapter 2.1 \\& 2.2 of Das, 2000a for detail solution scheme and flow geometry), we get the {\\it combined} flow topologies (Fig 2. of DC99) where the value of $R_{\\dot M}$ along with {\\it all} other outflow parameters can be exactly computed only from the knowledge of the {\\it minimum number of} accretion parameters namely specific energy (${\\cal E}$), specific angular momentum ($\\lambda$), accretion rate (scaled in the unit of Eddington rate) (${\\dot M}_{in}$) and adiabatic indices ($\\gamma$) of the flow. Also the dependence of  $R_{\\dot M}$ on {\\it all possible} flow parameters has been investigated self-consistently (see Das, 1998 and DC99 for detail). ",
        "conclusions": ""
    },
    "0212/astro-ph0212360_arXiv.txt": {
        "abstract": "We show that by combining measurements of the temperature and polarization anisotropies of the Cosmic Microwave Background (CMB), future experiments will tightly constrain the expansion rate of the universe during recombination. A change in the expansion rate modifies the way in which the recombination of hydrogen proceeds, altering the shape of the acoustic peaks and the level of CMB polarization. The proposed  test is similar in spirit to the examination of abundances of light elements produced during Big Bang Nucleosynthesis and it constitutes a way to study possible departures from standard recombination. For simplicity we parametrize the change in the Friedmann equation by changing the gravitational constant $G$. The main effect on the temperature power spectrum is a change in the degree of damping of the acoustic peaks on small angular scales. The effect can be compensated by a change in the shape of the primordial power spectrum. We show that this degeneracy between the expansion rate and the primordial spectrum can be broken by measuring CMB polarization. In particular we show that the MAP satellite could obtain a constraint for the expansion rate $H$ during recombination of $\\delta H/H \\simeq 0.09$ or $\\delta G/G \\simeq 0.18$ after observing for four years, whereas Planck could obtain $\\delta H/H \\leq 0.014$ or $\\delta G/G \\leq 0.028$ within two years, even after allowing for further freedom in the shape of the power spectrum of primordial fluctuations. ",
        "introduction": "}  The parameters of our cosmological model will be determined with great accuracy by upcoming data from Cosmic Microwave Background experiments, galaxy surveys, weak lensing surveys, Lyman alpha forest studies, and other observations. With the new data it will be possible to perform a number of consistency checks that will strengthen our confidence in the underlying model. Some of these consistency checks have already been performed with existing data.  Recent analysis of the CMB data have resulted in constraints on the baryon density $\\Omega_b h^2$ that are in excellent agreement with its determination based on the study of the primordial abundances of light elements (e.g. \\cite{Wang,Efsthathiou}). The combination of CMB data with local measures of the Hubble constant \\cite{keyproject} and measures of the local strength of galaxy clustering result in a determination of the cosmological constant that is in good agreement with results from the study of the luminosity of distant supernovae (e.g. \\cite{Riess,revpar}). Recently also joint BBN and CMB constraints on different dark energy models have been discussed in \\cite{Kneller:2002zh}.  One of the assumptions of the cosmological model that has been hard to test is the explicit validity of the Friedmann equation -- the relation between the expansion rate of the universe and its matter content. The difficulty lies in finding an epoch in the evolution of the universe during which both the energy density and the expansion rate can be determined independently. The two obvious candidates are the present time and Big Bang Nucleosynthesis (BBN).  Precise measures of both the expansion rate and the matter density in the local universe are difficult to accomplish and suffer from various systematic problems. At present the expansion rate has been measured with errors on the order of $10\\%$ \\cite{keyproject}. Direct determinations of the matter density however are more uncertain. It is fair to say that the general conclusion from these studies is that the Friedmann equation only holds true if either a cosmological constant or a curvature term is added. This is because direct determinations of the present matter density almost always point to $\\Omega_m < 1$ (e.g. see \\cite{revpar}).  Neither the cosmological constant nor the curvature scale can be constrained independently, that is without going through the Friedmann equation, so at best we can say that the Friedmann equation has not been tested accurately at the present epoch. A more radical interpretation would be that the Friedmann equation has been tested but that the test has failed. Although we do not support this interpretation, at present some infrared modification of gravity cannot be ruled out observationally. Such modifications are being explored for example as ways to solve the cosmological constant problem (see for example \\cite{Arkani-Hamed}) or to explain the accelerated expansion inferred from the luminosity distance to high redshift supernovae without resorting in a cosmological constant or a quintessence field (e.g. see \\cite{Deffayet:2001pu,Deffayet2}).  During the epoch of Big Bang Nucleosynthesis (BBN) the situation is more fortunate. The energy density is dominated by radiation which we think we can estimate accurately. On the other hand the expansion rate affects the freeze-out abundances of light elements, so that a precision test of the Friedmann equation can be performed.  The standard procedure is to constrain the number of relativistic degrees of freedom, $g_*$, which can be translated into a limit on the number of neutrino species. A lot of progress has been made in determining the primordial abundances. Recently the deuterium abundance in hydrogen clouds at high redshift was accurately determined \\cite{O'Meara:2000dh,Burles:1997ez,Burles:1997fa}. Building on a prior of $N_\\nu \\geq 3$ these data have been exploited to enforce an upper limit of 3.2 at $2\\sigma$ for the number of neutrino species, based purely on BBN considerations \\cite{Burles:1999zt}.  Progress has also been made in appreciating the systematical uncertainties that impair the determination of primordial $^4He$ (see e.g. \\cite{Olive:2000qm,Peimbert:2002ks}). Built on the safely established abundance ranges for Deuterium, Helium and Lithium, it can be shown that the uncertainty in the number of relativistic degrees of freedom during BBN is around 9 \\% (68\\% C.L.)\\cite{Olive:1998vj}. This constraint can equally well be phrased as a constraint on the validity of the Friedmann equation: during Nucleosynthesis the ratio of the squared expansion rate to the energy density can depart by only $9 \\%$ from what is predicted by the Friedmann equation. Constraints on the expansion history during BBN have also been established in \\cite{Carroll:2001bv}.    In this paper we propose using the anisotropies in the CMB to perform a test similar to the one that has been done using BBN.  We will show that such a test can ultimately constrain the validity of the Friedmann equation during recombination more accurately than what has so far been reached within Nucleosynthesis, albeit in a more model dependent way.  During recombination, the energy density is dominated by the density of non-relativistic matter which we cannot estimate directly. However the dark matter energy density enters in two different ways and one can exploit this to simultaneously determine the dark matter density and the expansion rate during recombination.  The ratio of matter to radiation energy density sets the redshift of matter radiation equality. At that time the expansion rate changes from a scaling as $t^{1/2}$ to $t^{2/3}$. Perturbation modes of the photon-baryon fluid that entered the horizon during the radiation dominated era behave differently than those that entered during the matter dominated era. Modes that entered during radiation domination provided the dominant contribution to the total density perturbation that generated the gravitational potential. On the contrary modes that entered the horizon during matter domination, were sub-dominant in their contribution to the total density perturbation which was dominated by the dark matter fluctuations. The gravitational potential acts as a source for perturbations in the photon baryon fluid. As a result, small scale modes that entered the horizon in the radiation era go through a sort of feedback loop that increases their amplitude as they cross the horizon (for a review of CMB physics see for example \\cite{DodlesonHu}). The anisotropy power spectrum is very sensitive to the redshift of matter radiation equality and thus the CMB should very accurately determine the ratio of dark matter to radiation energy density, i.e. the parameter $\\Omega_m h^2$.  The dark matter density dominates over other energy components in the Friedmann equation during recombination. In the standard scenario, it sets both the redshift at which matter and radiation become equal and the rate of expansion during recombination through the Friedmann equation. In this paper we break the link between energy density and expansion rate by introducing a free parameter to modify the Friedmann equation. We investigate how well such a parameter can be constrained.  From a pragmatic perspective our study can be regarded as the investigation of a particular departure from standard recombination. Different variations have been studied in the literature. For instance the possibility that energetic sources of Ly-$\\alpha$-photons could be present during recombination to delay it was considered in \\cite{Peebles:200l}. Also, the possiblility of a time variation of the fine structure constant was investigated \\cite{Kaplinghat:1998ry,Landau:2000dd,Hannestad:2000ys}. In \\cite{Uzan:2002vq} the effects of a time dependence of the gravitational constant have been outlined. The conclusion of these investigations and ours is that with future CMB data departures from standard recombination will be severely constrained and perhaps modifications that point to interesting new physics could be discovered.  In section \\ref{lambda}, we will introduce our model and in section \\ref{analysis} we will investigate the constraints that can be set with currently available data and forecast what future CMB experiments might be able to achieve. We will conclude in section \\ref{discussion} with discussion. ",
        "conclusions": "}   We have made the expansion rate of the universe a free parameter in a likelihood analysis within an eight-dimensional cosmological parameter space. For simplicity we assumed that the gravitational constant is changed by a factor $\\lambda$.  We showed that an increase of $\\lambda$ leads to a wider visibility function which in turn increases the damping of anisotropies on small scales and increases the level of large scale polarization.  We calculated the constraints that current CMB data can impose on the expansion rate.  The constraints that can be imposed on the parameter $\\lambda$ using the information from the damping tail are severely weakened by our lack of knowledge about the shape of the primordial power spectrum. We showed that measuring polarization helps to break this degeneracy.  Current data can only constrain $\\lambda$ to about 47\\% at 1 $\\sigma$.  We showed that MAP could obtain 9\\% error bars for $\\lambda$ while for Planck errorbars go down to 0.9\\%. We also explored the ultimate limit that could be achieved by a cosmic variance limited experiment measuring anisotropies up to $l=4000$ and found errors of under a percent in that case. Thus next generation experiments should be able to deliver very accurate constraints on the expansion rate of the universe during recombination.  We acknowledge that if the variation of the gravitational constant during recombination is taken seriously a model needs to be built where $G$ changes after recombination and converges towards the stable value observed in laboratory experiments today and where its current rate of change is less than the experimental bound $\\frac{\\dot{G}}{G} \\simeq 10^{-12} \\text{yr}^{-1}$\\cite{Uzan:2002vq,Guenther:1998}. If we introduce a scalar field to control the value of $G$ we would also have to require that this field does not lead to an unacceptably large fifth force and does not violate solar system constraints such as shifting the orbit of the moon through the Nordvedt effect \\cite{Will,Nordtvedt}.  The shift of $G$ after recombination will induce a change in the angular diameter distance to recombination, shifting the CMB power spectrum in $l$. We have shown that future experiments will be able to constrain the change of $G$ to a few percent due to its effect at recombination. As a result the induced shift in the peak positions would be small and could be interpreted as slightly different values of $\\Omega_m$ and/or $\\Omega_\\Lambda$, the two parameters that control the distance to the last scattering surface in conventional models. In the same way, any induced integrated Sachs-Wolfe (ISW) effect would be difficult to observe because of the cosmic variance limitation. If observed with a high enough accuracy it should not be identical with what is predicted by a simple cosmological constant model.  In this paper we have shown that future measurements of the CMB anisotropy will be able to extract information about the relation between the expansion rate and the energy density of the universe during recombination, because of its effect on the recombination history of hydrogen."
    },
    "0212/astro-ph0212010_arXiv.txt": {
        "abstract": "Two questions about the solar magnetic field might be answered together once their connection is identified. The first is important for  large scale dynamo theory: what prevents the magnetic backreaction forces from shutting down the dynamo cycle? The second question is: what determines the handedness of twist and writhe in magnetized coronal ejecta? Magnetic helicity conservation is important for answering both questions. Conservation implies that dynamo generation of large scale writhed structures is accompanied by the oppositely signed twist along these structures. The latter is associated with the backreaction force. We suggest that coronal mass ejections (CMEs) simultaneously liberate small scale twist and large scale writhe of opposite sign, helping to prevent the cycle from quenching and enabling a net magnetic flux change in each hemisphere. Observations and helicity spectrum measurements from a numerical simulation of a rising flux ribbon in the presence of rotation support this idea.  We show a new pictorial of dynamo flux generation that includes the backreaction and magnetic helicity conservation by characterizing the field as a 2-D ribbon rather than a 1-D line.  \\medskip  {{\\bf Subject Headings}: MHD--Sun: activity--Sun: magnetic fields MHD--turbulence; stars--magnetic fields; galaxies--magnetic fields; methods--numerical} ",
        "introduction": "The helical magnetic dynamo is the basis for a promising class of mechanisms to explain large scale magnetic fields observed in stars and galaxies \\cite{parker55,moffatt78,krause,parker93}. The basic ``$\\alpha-\\Omega$'' dynamo is the specific version most relevant for strongly sheared rotators (Fig.~1). Interface $\\alpha-\\Omega$ dynamos (Parker 1993; Charbonneau \\& MacGregor 1996; Markiel \\& Thomas 1999) include the fact that, unlike for galaxies and disks, the dominant shear layer is beneath the dominant turbulent region.   Focusing on the simplest ``$\\alpha-\\Omega$'' picture (Fig.~1), consider an initially weak toroidal (=encircling the rotation axis) loop of the magnetic field embedded in the astrophysical plasma rotator. The magnetic field is coupled to the plasma so the field is carried or stretched in response the plasma motion. Now imagine, as in the sun, that there is an outwardly decreasing density gradient. Consider the motion of a rising turbulent swirl of gas, threaded by a magnetic field. Conservation of angular momentum dictates that the swirl will writhe oppositely to the underlying  rotation of the system. This means that the initially toroidal field threading the swirl gains a radial component. Statistically, rising swirls in the northern (southern) hemisphere writhe the field clockwise (counterclockwise). This is the ``$\\alpha$'' effect and is shown by the writhed loop of Fig.~1a for the northern hemisphere. Differential rotation at the base of the loop shears the radial field (the ``$\\Omega$''-effect). The bottom part of the loop amplifies the initial toroidal seed loop as shown in Fig.~1b, whilst the top part of the loop diffused away (the ``$\\beta$'' effect). In doing so, magnetic flux is amplified: the flux penetrating the tilted rectangular surface is zero in Fig.~1a, but finite in Fig.~1b.  This process is represented mathematically by averaging the magnetic induction equation over a local volume and breaking all quantities (velocity $\\bf U$, magnetic field $\\bf B$ in Alfv\\'en velocity units, and  normalized current density ${\\bf J}\\equiv {\\curl {\\bf B}}$) into their mean (indicated by an overbar) and fluctuating (indicated by lower case) components. The result is \\cite{moffatt78}: $\\partial_t\\meanBB= \\curl(\\alpha\\meanBB+{\\overline {\\bf U}}\\ts \\meanBB) +(\\beta+\\lambda) \\nabla^2\\meanBB, $ where $\\lambda$ is the microphysical diffusivity. The ${\\overline {\\bf U}}$ term incorporates the $\\Omega$-effect, the $\\beta$ term incorporates the turbulent diffusion (assuming constant $\\beta$) and the first term on the right incorporates the $\\alpha$-effect. In the kinematic theory \\cite{moffatt78} $\\alpha$ is given by $\\alpha =\\alpha_{0}= -(\\tau/3) {\\overline {\\bfu\\cdot\\curl\\bfu}}$. Here $\\tau$ is a turbulent damping time and ${\\overline {\\bfu\\cdot\\curl\\bfu}}$ is the kinetic helicity, which dictates the $\\alpha$-effect described physically above. Usually, $\\alpha$ and $\\beta$ are prescribed as input parameters.  A long standing problem has been the absence of properly incorporating the (time-dependent) backreaction from the growing magnetic field on the driving turbulent motions. This stimulated criticisms of mean-field dynamos \\cite{piddington,cattaneohughes} and motivated interface dynamo models \\cite{parker93}. But the backreaction is now much better understood. Steady-state studies of $\\alpha$-quenching \\cite{cattaneohughes} apply only at fully saturated dynamo regimes, not at early times, when the backreaction is just beginning to be important. There the growth is fast, and most relevant for astrophysics. Demonstrating this  requires including the time-evolution of the turbulent velocity, subject to magnetic forces. Carrying this out formally \\cite{bf02} and using a closure in which triple correlations act as a damping term, amounts to replacing $\\alpha =\\alpha_{0}$  by $\\alpha=\\alpha_0+ (\\tau/3){\\overline {\\bfb\\cdot{\\curl \\bfb}}}$, where the second term is the backreaction. It arises from $({\\bf \\curl \\bfb}) \\times \\bbB$, the force associated with the action of the small scale current and the large scale field. This residual form of $\\alpha$ has long been thought \\cite{pfl} to be the real driver of the helical dynamo and has been employed in attempts to understand its quenching \\cite{zeldovich83,kleeorin82,fb02,bb02} From the form of $\\alpha$, it is evident that a large current helicity can offset the kinetic helicity and quench the dynamo \\cite{fb02,bb02}.  In section 2 we summarize the successful backreaction theory, and show how it predicts ejection of twist and writhe of opposite sign. In section 3 we give a new pictorial of dynamo action that includes magnetic helicity conservation, and discuss a simulation of a rising flux ribbon. Observational implications are discussed in section 4, and we conclude in section 5. ",
        "conclusions": "We have suggested that a new understanding of how helical dynamos conserve magnetic helicity may help resolve several mysteries of the solar magnetic field. We have shown that large and small scale helicities of approximately equal magnitude should be ejected into the solar corona as part of the sustenance of the solar cycle. The large scale helicity  corresponds to the writhe of a prominence, whilst the small scale helicity corresponds to the twist along the prominence. We emphasize the importance of simultaneously detecting large and small scale contributions to the losses of helical magnetic fields in pre-CME sigmoid structures at the solar surface. The observations seem to be roughly consistent with our simple picture at present but more studies and detailed modeling will be needed to test these ideas. Our Fig.~1 illustrates the basic concepts through a new pictorial representation of the mean field dynamo that includes magnetic helicity conservation and the backreaction. The implications are also relevant for large scale dynamos in other stars and disks in astrophysics.                     \\noindent EB acknowledges DOE grant DE-FG02-00ER54600, and the Dept.\\ of Astrophys.\\ Sciences at Princeton  for hospitality during a sabbatical. We acknowledge  PPARC supported supercomputers at Leicester and St Andrews, and the Odense Beowulf cluster  is acknowledged."
    },
    "0212/astro-ph0212516.txt": {
        "abstract": "We present new gas flow models for the Milky Way inside the solar circle. We use SPH simulations in gravitational potentials determined from the NIR luminosity distribution of the bulge and disk, assuming constant NIR mass-to-light ratio, with an outer halo added in some cases. The luminosity models are based on the {\\it COBE/DIRBE} maps and on clump giant star counts in several bulge fields, and include a spiral arm model for the disk.  Gas flows in models which include massive spiral arms clearly match the observed $^{12}$CO \\lvplot\\ better than if the potential does not include spiral structure. Furthermore, models in which the luminous mass distribution and the gravitational potential of the Milky Way have four spiral arms are better fits to the observed \\lvplot\\ than two-armed models.  Besides single pattern speed models we investigate models with separate pattern speeds for the bar and spiral arms.  The most important difference is that in the latter case the gas spiral arms go through the bar corotation region, keeping the gas aligned with the arms there. In the ($l,v$) plot this results in characteristic regions which appear to be nearly void of gas.  In single pattern speed models these regions are filled with gas because the spiral arms dissolve in the bar corotation region.  Comparing with the $^{12}$CO data we find evidence for separate pattern speeds in the Milky Way.  From a series of models the preferred range for the bar pattern speed is $\\Omega_p=60\\pm 5 \\Gyr^{-1}$, corresponding to corotation at $3.4\\pm0.3 \\kpc$. The spiral pattern speed is less well constrained, but our preferred value is $\\Omega_{sp}\\approx 20 \\Gyr^{-1}$. A further series of gas models is computed for different bar angles, using separately determined luminosity models and gravitational potentials in each case. We find acceptable gas models for $20\\deg\\lta\\phibar \\lta25\\deg$. The model with ($\\phibar=20\\deg$, $\\Omega_p=60\\Gyr^{-1}$, $\\Omega_{sp}=20 \\Gyr^{-1}$) gives an excellent fit to the spiral arm ridges in the observed ($l,v$) plot. ",
        "introduction": "\\label{s1} Observations of cold gas in the Milky Way [MW] have contributed substantially to our understanding of MW structure. No other tracer is observed in as large a part of the MW as are gas clouds. Longitude-velocity [lv] diagrams (Hartmann \\& Burton \\shortcite{HB97}, Dame et al.\\ \\shortcite{Dame01}) show the distribution of gas velocities as a function of galactic longitude $l$, integrated over some range in latitude $b$.  By observing the MW in different spectral lines, this gas can be traced at substantially different densities. The largest absolute velocity as a function of $l$ defines the terminal velocity curve (\\tvc). In an axisymmetric galaxy, the gas at these velocities is found at the ``tangent point'' where the line of sight is tangential to a circle around the Galactic centre. From this the rotation curve can be determined. However, due to the bar and spiral perturbations in the MW potential, the gas has substantial non-circular velocities, which are most evident in the central $10\\deg-20\\deg$, due to the bar, but also as ``bumps'' in the \\tvc\\ where spiral arm tangents perturb the gas flow by $\\sim 10$-$20 \\kms$. At sub-\\tvc\\ velocities, crowding in both position and in velocity produces ridge-like structures in the \\lvplot. The Galactic spiral arms are visible as straight or curved such ridges.  A number of attempts have been made to model these observations.  One group is formed by analytic models of spiral structure.  The first exhaustive analytical formulation of a spiral arm theory was developed by Lin \\& Shu \\cite{Lin64} and applied to the MW by Lin et al.\\ \\cite{Lin69}. They proposed a two-armed model with pitch angle $-6\\deg$ and a pattern speed for the spiral structure $\\Omega_{sp} \\approx 13.5 {\\rm km/s/kpc}$. Amaral \\& L\\'epine \\cite{AmLep97} fitted the rotation curve of the MW to an analytic mass model and found a self-consistent solution with a combined two- and four-armed spiral structure. In L\\'epine et al.\\ \\cite{Lep01} they extended this model to allow for a phase difference between the two-armed and the four-armed spiral pattern.  The second group, numerical simulations of the Galactic gas flow, also have a long tradition.  A recent example is the smoothed particles hydrodynamics (SPH) models of Fux \\cite{Fux99}, who evolved a gas disk inside a self-consistent $N$-body model scaled to the {\\it COBE/DIRBE} K-band map of the MW and the radial velocity dispersion of M giants in Baade's window. The resulting gas flow was transient, but at specific times resembled closely a number of observed arms and clumps in the bar region.  Weiner \\& Sellwood \\cite{Weiner99} compared predictions from fluid dynamical simulations for analytic mass densities with the observed outer velocity contours of the HI $(l,v)$ diagram to constrain the pattern speed and bar angle.  Englmaier \\& Gerhard \\shortcite{EG99}  [hereafter, Paper I] computed gas flows in the gravitational potential of the NIR luminosity distribution of Binney, Gerhard \\& Spergel \\cite{BGS97}, assuming constant NIR mass-to-light ratio ($M/L$). Their best SPH gas flow models reproduced quantitatively a number of observed gas flow features, including the positions of the five main spiral arm tangents at $|l|\\leq 60\\deg$ and much of the terminal velocity curve.  An important feature of all these gas flow models is the pattern speed of the non-axisymmetric component. Englmaier \\& Gerhard found a best pattern speed for the bar $\\Omega_p\\approx 60 {\\rm km/s/kpc}$.  Weiner \\& Sellwood derived a bar pattern speed $\\approx 42 {\\rm km/s/kpc}$, whereas Fux determined $\\approx 50 {\\rm km/s/kpc}$ from his models.  Dehnen \\cite{Dehn00} used resonant features in the Hipparcos stellar velocity distribution to argue that the Sun is located just outside the OLR of the exciting quadrupole perturbation, giving a pattern speed of $\\approx 51 {\\rm km/s/kpc}$ for solar constants $R_0=8\\kpc$ and $v_0=220\\kms$. For their {\\it spiral structure} model, Amaral \\& L\\'epine \\cite{AmLep97} and L\\'epine et al.\\ \\cite{Lep01} found $\\Omega_{sp}\\approx 20\\, \\Gyr^{-1}-35\\, \\Gyr^{-1}$.  Fern\\'andez \\etal\\ \\cite{Fer01} obtained a somewhat higher $\\Omega_{sp}\\approx 30 {\\rm km/s/kpc}$ from Hipparcos data for OB stars and Cepheids. Debattista et al.\\ \\cite{Deba02} used the Tremaine-Weinberg method on a sample of intermediate age to 8 Gyr old OH/IR-stars in the inner Galactic disk.  They found a pattern speed of $59 \\pm 5 \\pm 10 $ (systematic) km/s/kpc, which may be driven by the bar in the center of the MW.   Thus the bar and spiral arms in the MW may not rotate with the same pattern speed.  For a fast bar, a single $\\Omega_p$ would imply that the spiral arms are entirely {\\sl outside} their corotation radius. Observations of external galaxies (see, e.g., the Hubble Atlas of Galaxies, Sandage \\shortcite{Sandage61}) suggest that galaxies exist with dust lanes on the inner (concave) edges of their spiral arms. For a trailing spiral pattern, these arms would be {\\it inside} their corotation radius.  A lower pattern speed for the spiral structure than for the bar would remove this discrepancy.  Indeed, Sellwood \\& Sparke \\cite{Sell88} showed evidence for multiple pattern speeds in their $N$-body simulations.  Rautiainen \\& Salo \\cite{Rau99} analysed two-dimensional $N$-body simulations, some of them with a massless, dissipative gas component added.  They confirmed the possibility of multiple pattern speeds in self-consistent N-body models of barred galaxies, and found a number of possible configurations. These included models with corotating bar and spirals, as well as models with different pattern speeds. Some of the models in the latter group show evidence for a non-linear mode-coupling (Tagger \\etal\\  \\shortcite{Tagger87}) between bar and spiral pattern, but others show no such evidence.  In some of their models there exist separate inner spirals corotating with the bar, and outer spirals which rotate with their own, lower pattern speed.  It is therefore tempting to analyse gas flow models of the MW with multiple pattern speeds.  What is the morphology of the MW spiral arms?  Most authors infer four spiral arms from tracers which directly or indirectly measure the gas density, such as molecular clouds, HII regions, pulsars, and the galactic magnetic field (Georgelin \\& Georgelin \\shortcite{gg76}, Sanders et al.\\ \\shortcite{Sanders85}, Caswell \\& Haynes \\shortcite{CasHay87}, Grabelsky \\& et al.\\ \\shortcite{Grab88}, Taylor \\& Cordes \\shortcite{TayCor93}, Vall\\'ee \\shortcite{Val95}; however, Bash \\cite{Bash81} infers a two-armed pattern from the same HII-data as used by Georgelin \\& Georgelin). The problem is that all spiral arm parameters other than the tangent point directions (e.g., Paper I) require distance information. It is also not clear whether all spiral arms seen in the MW gas are present in the old disk. Ortiz \\& L\\'epine \\cite{Ort93} constructed a four-armed model which reproduces their star counts in the near infrared.  Drimmel \\cite{Drimmel00} preferred a two-armed structure from {\\it COBE/DIRBE} K-band data, but a four-armed structure for the dust distribution seen in the $240\\mum$ data.  Drimmel \\& Spergel \\cite{Drimmel01} project a luminosity model through the $240\\mum$ dust model, to compare with the NIR J and K band {\\it COBE/DIRBE} data.  Their best model for the stellar distribution is four-armed, but dominated by two spiral arms.  Drimmel \\& Spergel conclude that, if there are four arms in the K-band luminosity distribution, the Sag-Car arm is of reduced strength (by a factor of $2.5$).  In this paper we investigate the dynamical effects of the Galactic bar and spiral arms on the gas flow in the Milky Way.  We investigate the possibility of different pattern speeds for bar and spiral arms, and the consequences this would have on the observed \\lvplots.  Our mass models for the inner Galaxy are based on the NIR luminosity density models of Bissantz \\& Gerhard \\cite{BissGerh02} [hereafter, Paper II] which include spiral structure.  We use SPH simulations to determine the gas flow in the MW, extending the work presented in Paper I, where eightfold symmetric mass models were used.  This paper is organised as follows.  In Section \\ref{s2} we describe the luminosity models, methods, and observational data used in this work.  In Section \\ref{s3} we describe our best gas model for the observed $^{12}$CO \\lvplot. Then we compare models with different pattern speeds (Section \\ref{s4}), bar angles and spiral arm morphology (Section \\ref{s5}) to constrain these parameters, and finally give our conclusions in Section \\ref{s7}. ",
        "conclusions": "\\label{s7} We have used new gas flow models to investigate the dynamics of the Milky Way (MW) Galaxy from observed \\lvplots.  Steady-state gas flows in rotating, point-symmetric gravitational potentials for the Galactic bar and disk were determined with SPH simulations. The potentials were derived from non-parametric estimates of the spatial near-infrared luminosity density (Bissantz \\& Gerhard \\shortcite{BissGerh02}), based on the de-reddened {\\it COBE/DIRBE} L-band map of Spergel et al.\\ \\cite{Spergel96}, but also incorporating clump giant star count data from Stanek \\etal\\ \\shortcite{Stanek94}, \\shortcite{Stanek97}.  The luminosity models contain a spiral arm model for the disk, and were converted to mass models assuming constant mass-to-light ratio in the inner MW.  Our best gas flow model gives a very good fit to the Galactic terminal velocity curve for $|l|>15\\deg$, and to the spiral arm ridges in the observed CO \\lvplot. This has enabled us to investigate a number of dynamically important parameters such as the bar and spiral arm pattern speeds, the multiplicity of the spiral structure in the potential, and the bar angle. The main results from this study are as follows.   1) In gas flow models with separate pattern speeds $\\Omega_p$ for the bulge/bar and $\\Omega_{sp}$ for the spiral pattern, the spiral arms go through the bar corotation region. Thus \\lvplots\\ of such models show well-defined spiral arm shocks (ridges) through corotation, next to areas which appear to be nearly void of gas.  By contrast, in single pattern speed models the spiral arms dissolve in the bar corotation region, so that the gas fills this region of the \\lvplot\\ approximately evenly, and no voids exist.   2) Similar voids are visible in the observed $^{12}$CO \\lvplot. From a comparison with model \\lvplots\\ with different $\\Omega_{sp}$ but similar $\\Omega_p$ we find evidence for separate pattern speeds in the Milky Way. The existence of self-consistent models with separate bar and spiral arm pattern speeds was demonstrated by Rautiainen \\& Salo \\shortcite{Rau99} in a study of two-dimensional N-body simulations, some of which included a massless, dissipative gas component. Models with a growing bar amplitude also support spiral arms in the corotation region  (Thielheim \\& Wolff \\shortcite{TW82}); however, it is likely that when self-gravity is included and the spiral arm amplitude becomes non-linear, the growing spiral pattern will again develop an independent pattern speed.  3) From a series of models the preferred range for the bar pattern speed in the MW is $\\Omega_p=60\\pm 5 \\Gyr^{-1}$, corresponding to corotation at $3.4\\pm 0.3 \\kpc$.  This agrees well with previous pattern speed determinations by Englmaier \\& Gerhard \\cite{EG99}, Dehnen \\cite{Dehn00}, and Debattista \\etal\\ \\cite{Deba02}. The bar pattern speed is well constrained because it influences not only the inner spiral structure, but also the position of two outer spiral arms in the lv-plot.  Models with $\\Omega_p=50 \\Gyr^{-1}$ and $\\Omega_p=70 \\Gyr^{-1}$ are inferior.  The spiral arm pattern speed is less well constrained. Our preferred value is $\\Omega_{sp}\\approx 20 \\Gyr^{-1}$, but models with larger $\\Omega_{sp} \\lt \\Omega_p$ give only marginally inferior fits to the observed \\lvplot.   4) Gas flows in models which include massive spiral arms clearly fit the observed $^{12}$CO ($l,v$) plot better than if the potential does not include spiral structure. Furthermore, comparing models with two and four arms in the gravitational potential, we found that only models with four massive arms reproduce the Galactic \\lvplot, while gas flows in two-armed potentials do not resemble the spiral arm pattern of the Milky Way.  In Galactic models with four-armed potentials and separate spiral arm pattern speed, the gas flow has two pairs of inner arms which rotate with the bar (lateral, and corresponding to the 3 kpc arm), and four outer spiral arms which exhibit a complicated, time-dependent back--and--forth oscillation in the bar frame. The outer and inner spiral structures are connected by a time-dependent transition region around bar corotation.   5) From a further series of gas models computed for different bar angles, using separately determined luminosity models and gravitational potentials as in Bissantz \\& Gerhard \\shortcite{BissGerh02}, we found a range of acceptable bar angles $20\\deg\\lta\\phibar \\lta25\\deg$. The models for $\\phibar= 15\\deg$ and $30\\deg$ are clearly inferior, which is mainly due to differences in the inferred gravitational potential.  The model with ($\\phibar=20\\deg$, $\\Omega_p=60\\Gyr^{-1}$, $\\Omega_{sp}=20 \\Gyr^{-1}$) gives an excellent fit to the Galactic terminal velocity curve for $10\\deg \\leq |l| \\leq 50\\deg$, and to the gaps and spiral arm ridges in the observed CO \\lvplot. There are still discrepancies in the bar corotation region where the potential has uncertainties: the 3 kpc arm has too low non-circular velocity, and its counterarm is missing in the data. In the bulge region, closed orbits reproduce the \\tvc\\ well, while the gas model has lower velocities. This may be a resolution problem in the SPH model, but could in part also be due to uncertainties in the potential which influence the orbit shapes.   6) The 3 kpc arm non-circular velocities can be reproduced by a model in which we artifically increased the $m\\geq 2$-multipoles of the spiral potential component by $1.5$ while keeping all other dynamical parameters fixed. This is well within the uncertainties. Guided by a number of asymmetries in the observed Milky Way gas distribution, we also investigated potentials which are no longer point-symmetric. Some of these models improved the fit to the 3kpc-arm and its counter-arm. Although we did not find a model which at the same time reproduces the entire \\lvplot\\ as well as our standard model, these models suggest that such asymmetries may be important for better understanding the gas flow in the inner Milky Way."
    },
    "0212/astro-ph0212226_arXiv.txt": {
        "abstract": "\\noindent{\\it Spectroscopic analyses with the intention of the interpretation of the UV-spectra of the brightest stars as individuals -- supernovae -- or as components of star-forming regions -- massive O~stars -- provide a powerful tool with great astrophysical potential for the determination of extragalactic distances and of the chemical composition of star-forming galaxies even at high redshifts.  The perspectives of already initiated work with the new generation of tools for quantitative UV-spectroscopy of Hot Stars that have been developed during the last two decades are presented and the status of the continuing effort to construct corresponding models for Hot Star atmospheres is reviewed.  Since the physics of the atmospheres of Hot Stars are strongly affected by velocity expansion dominating the spectra at all wavelength ranges, hydrodynamic model atmospheres for O-type stars and explosion models for Supernovae of Type~Ia are necessary as basis for the synthesis and analysis of the spectra. It is shown that stellar parameters, abundances, and stellar wind properties can be determined by the methods of spectral diagnostics already developed. Additionally, it will be demonstrated that models and synthetic spectra of Type~Ia Supernovae of required quality are already available. These will make it possible to tackle the question of whether Supernovae~Ia are standard candles in a cosmological sense, confirming or disproving that the current SN-luminosity distances indicate accelerated expansion of the universe.  In detail we discuss applications of the diagnostic techniques by example of two of the most luminous O~supergiants in the Galaxy and a standard Supernova of Type~Ia. Furthermore, it is demonstrated that the spectral energy distributions provided by state-of-the-art models of massive O~stars lead to considerably better agreement with observations if used for the analysis of H\\,{\\scriptsize II}~regions. Thus, an excellent way of determining extragalactic abundances and population histories is offered.  Moreover, the importance of Hot Stars in a broad astrophysical context will be discussed. As they dominate the physical conditions of their local environments and the life cycle of gas and dust of their host galaxies, special emphasis will be given to the corresponding diagnostic perspectives. Beyond that, the relevance of Hot Stars to cosmological issues will be considered.} ",
        "introduction": "It is well known that Hot Stars are not a single group of objects but comprise sub-groups of objects in different parts of the HR~diagram and at different evolutionary stages. The most important sub-groups are massive O/B stars, Central Stars of Planetary Nebulae, and Supernovae of Type~Ia and~II. All these sub-groups have in common that they are characterized by high radiation energy densities and expanding atmospheres, and due to this the state of the outermost parts of these objects is characterized by non-equilibrium thermodynamics. In order to cover the best-known fundamental stages of the evolution of Hot Stars in sufficient depth this review will be restricted to the discussion of O~stars and Supernovae of Type~Ia; we will not discuss objects like Wolf-Rayet stars, Luminous Blue Variables, Be-stars, Supernovae of Type~II, and others.  Furthermore, Hot Stars play an important role in a broad astrophysical context. This implies that a complete review covering all aspects in theory and observation is not only impossible, but also beyond the scope of this review. Thus, this review will focus on a special part of the overall topic with the intent to concentrate on just one subject; the subject chosen is UV spectral diagnostics. It will be shown, however, that this subject has important implications for astrophysical topics which are presently regarded as being ``trendy''.  But first of all we have to clarify what UV spectral diagnostics means. This is best illustrated by the really old-fashioned (1977, Morton and Underhill) UV~spectrum of one of the brightest massive O~stars, the O4\\,I(f) supergiant $\\zeta$~Puppis. As can be seen in Figure~1, expanding atmospheres have a pronounced effect on the emergent spectra of hot stars -- especially in the UV-part. The signatures of outflow are clearly recognized by the blue-shifted absorption and red-shifted emission in the form of the well-known P Cygni profiles. It is quite obvious that these kind of spectra contain information not only about stellar and wind parameters, but also about abundances. Thus, in principle, all fundamental parameters of a hot star can be deduced from a comparison of observed and synthetic spectra.  Although spectra of the quality shown in Figure~1 have been available for more than 25 years, most of the work done to date has concentrated on qualitative results and arguments. In view of the effort put into the development of modern -- state-of-the-art -- instruments, it is certainly not sufficient to restrict the analyses to simple line identifications and qualitative estimates of the physical properties. The primary objective must be to extract the complete physical stellar information from these spectra. Such diagnostic issues principally have been made possible by the superb quality and spectral resolution of the spectra available.  For this objective, the key is to produce realistic synthetic UV spectra for Hot Stars. Such powerful tools, however, are still in development and not yet widely used. They offer the opportunity to determine the stellar parameters, the abundances of the light elements -- He, C, N, O, Si -- and of the heavy elements like Fe and Ni, quantitatively, as is indicated not only by IUE, but also more recent HST, ORFEUS, and FUSE observations of hot stars in the Galaxy and Local Group galaxies. All these observations show that the spectra in the UV spectral range are dominated by a dense forest of slightly wind-affected pseudo-photospheric metal absorption lines\\linebreak \\myfigure{\\includegraphics[width=11.0cm]{r1606_fig01.eps}} Figure 1: {\\small Merged spectrum of Copernicus and IUE UV high-resolution observations of the O4\\,I(f) supergiant $\\zeta$ Puppis (900--1500~\\AA: Morton and Underhill~1977; 1500--1800~\\AA: Walborn et\\,al.\\ 1985). The most important wind lines of the light elements are identified and marked. Also marked are the large number of wind-contaminated lines of the iron group elements (e.g., \\FeV) which are especially present between 1250 and 1500~\\AA. (Figure from Pauldrach et\\,al.\\ 1994b).}\\\\  \\noindent overlaid by broad P~Cygni line profiles of strong, mostly resonance lines, formed in different parts of the expanding atmosphere. Thus, the obvious objective is to investigate the importance of these lines with respect to the structure of the expanding atmospheres that are characterized by the strength and the velocity of the outflow, and through which the shape of the spectral lines is mainly determined.  Over the last 30 years it turned out that the achievement of this objective, which will also be the primary aim of this paper to review, remains a difficult task.  Before we discuss in detail the status quo of the diagnostic tool required (Section~6), we will first examine whether such a tool has been made obsolete by the general development in astrophysics or whether it is still relevant to current astronomical research.  For this purpose we first discuss the diagnostic perspectives of galaxies with pronounced current star formation. Due to the impact of massive stars on their environment the physics underlying the spectral appearance of starburst galaxies are rooted in the atmospheric expansion of massive O~stars which dominate the UV wavelength range in star-forming galaxies. Therefore, the UV-spectral features of massive O~stars can be used as tracers of age and chemical composition of starburst galaxies even at high redshift (Section~2).  With respect to the present cosmological question of the reionization of the universe -- which appeared to have happened at a redshift of about $z \\sim 6$ -- the ionization efficiency of a top-heavy Initial Mass Function for the first generations of stars is discussed (Section~3).  Starting from the impact of massive stars on their environment it is demonstrated that the spectral energy distributions provided by state-of-the-art models of massive O~stars lead to considerable improvements if used for the analysis of \\HII~regions. Thus, the corresponding methods for determining extragalactic abundances and population histories are promising (Section~4).  Regarding diagnostic issues, the role of Supernovae of Type~Ia as distance indicators is discussed. The context of this discussion concerns the current and rather surprising surprising result that distant SNe~Ia appear fainter than standard candles in an empty Friedmann model of the universe (Section~5).  Finally, we discuss applications of the diagnostic techniques by example of two of the most luminous O~supergiants in the Galaxy; additionally, basic steps towards realistic synthetic spectra for Supernovae of Type~Ia are presented (Section~7). ",
        "conclusions": "The \\textbf{need} for \\textbf{detailed atmospheric models} of \\textbf{Hot Stars} has been motivated in depth and a \\textbf{diagnostic tool} with great astrophysical potential has been presented.  It has been shown that the models of the {\\it new generation\\/} are realistic. They are {\\it realistic\\/} with regard to a {\\it quantitative spectral UV analysis\\/} calculated, for the case of O~stars, along with consistent dynamics, which, in principle, allows to read off the stellar parameters by comparing an observed UV~spectrum to a suitable synthetic spectrum. The new generation of models has reached a degree of sophistication that makes such a procedure practicable.  The astronomical perspectives are enormous. Of course, there is still a large amount of hard and careful astronomical work that needs to be done. The diagnostic tool for expanding model atmospheres is in our hands, but this is (as usual) just the starting point; further elaboration, refinements, and modifications are required before the results are quantitatively completely reliable.  But this is not the handicap for the future. The present handicap are future observations. What is most urgently needed is the \\textbf{Next Generation Space Telescope}.  With respect to the topics discussed in this review this telescope is needed for a variety of reasons:  To tackle the question whether \\textbf{Supernovae~Ia} are standard candles in a cosmological sense realistic models and synthetic spectra of Type~Ia Supernovae require {\\it spectral observations of objects at mid-redshift\\/} and sufficient resolution. The context of this question is the current surprising result that distant SNe~Ia at intermediate redshift appear fainter than standard candles in an empty Friedmann model. Consequently, the current SN-luminosity distances indicate an accelerated expansion of the universe.  In order to be able to make quantitative predictions about the influence of very massive, extremely metal-poor \\textbf{Population~III} stars on their galactic and intergalactic environment one primarily needs observations which can be compared to the predicted flux spectra that are already available for zero metallicity and can, in principle, be produced for metallicities different from zero. Observations with the Next Generation Space Telescope (NGST) of distant stellar populations at high redshifts will give us the opportunity to deduce the primordial IMF, thus allowing a quantitative investigation of the ionization efficiency of a Top-heavy IMF via {\\it realistic spectral energy distributions\\/} of these very massive stars.  NGST observations are also important for determining extragalactic abundances and population histories of \\textbf{starburst galaxies} from an analysis of \\HII~regions. This task also requires energy distributions of time-evolving stellar clusters, which are calculated on the basis of a grid of \\textit{spectral energy distributions} of O and early B-type stars. A crucial point with respect to this is whether the spectral energy distributions of massive stars are already realistic enough to be used for diagnostic issues of \\HII~regions.  It has been shown in this review that the \\textit{ultimate test} is provided by a comparison of observed and synthetic UV~spectra of individual massive stars, since the ionization balance can be traced reliably through the strength and structure of the wind lines formed throughout the atmosphere.  Furthermore, NGST is important for directly exploiting the diagnostic perspectives of galaxies with pronounced current star formation. It has been demonstrated that massive stars dominate the UV wavelength range in \\textbf{star-forming galaxies}, and that therefore the UV-spectral features of massive O~stars can be used as tracers of age and chemical composition of starburst galaxies even at high redshift. This is in particular the case when the flux from these galaxies is amplified by gravitational lensing through foreground galaxy clusters; corresponding observations render the possibility of determining metallicities of starbursting galaxies in the early universe via {\\it realistic UV~spectra} (in the rest frame) of massive O~stars.  \\smallskip As a final remark it is noted that the solution method of stationary models for expanding atmospheres is in its present stage already regarded as a standard procedure towards a realistic description. Thus, together with an easy-to-use interface and an installation wizard, the program package \\textbf{WM-basic} has been made available to the community and can be downloaded from the author's home page.   \\subsection*"
    },
    "0212/astro-ph0212232_arXiv.txt": {
        "abstract": "{ We report on the results of a two epoch study of the low power Compact Symmetric Object 4C\\,31.04. Observations performed with the VLBA at 5 GHz in 1995 and 2000 have yielded images of this source at milliarcsecond angular resolution. A central core is detected, with bright compact hot spots and extended lobes on both sides. Model-fitting and other analysis of the data (brightness profile, difference map) clearly indicate that the source is expanding. We estimate the velocity of this expansion to be (0.085 $\\pm$ 0.016) mas/yr, i.e. (0.33 $\\pm$ 0.06) $h_{65}^{-1} \\ c$ in both hot spots. Assuming a constant expansion velocity, we estimate the kinematic age of the source at 550 yrs. We also study the spectral index using VLBA observations at 1.3 GHz and MERLIN at 22 GHz. The derived spectral age is 3000-5000 years in equipartition conditions. The two estimates are discussed and found to be in agreement, given present uncertainties. ",
        "introduction": "In recent years, considerable efforts have been devoted to the study of Compact Symmetric Objects (CSOs).  The name for this class of objects is derived from their morphology -- they have sub-galactic dimensions but the same symmetric, two-sided aspect as classical large scale FR\\,I and FR\\,II radio galaxies.  The CSOs also show other interesting properties in the radio: (1) their spectrum can peak at frequencies as high as a few GHz (hence their frequent classification as Gigahertz Peaked Spectrum -- see O'Dea (\\cite{ode98}), Polatidis \\& Conway (\\cite{pc}) and Polatidis et al. (\\cite{pol02}) for recent reviews), (2) the rate of H\\kern0.1em{\\sc i}\\ absorption is higher than in common radio galaxies (e.g. Van Gorkom et al. \\cite{jvg89}; Pihlstr\\\"om \\cite{pih01}), and (3) variability and polarization are almost non-existent (Fassnacht \\& Taylor \\cite{fas01}). At other wavelengths a comprehensive study of the properties of CSO host galaxies is still missing, although HST observations have suggested that in at least some cases the host galaxies are not completely relaxed (Perlman et al. \\cite{per01}).  Our current understanding of CSOs suggests that they are small by virtue of their youth (\\textit{youth scenario}) and not confined by an extremely dense medium (\\textit{frustration scenario}).  If they are young then CSOs represent the very first stage of an evolutionary sequence (Fanti et al. \\cite{fan95}; Readhead et al. \\cite{rea96a}, \\cite{rea96b}; Polatidis \\& Conway \\cite{pc}) in which a radio-quiet galaxy turns into the host of a large, kpc-scale radio loud object. It is therefore of great interest to (1) investigate evidence supporting and quantifying the age of CSOs and (2) relate their current properties to the mechanisms triggering the nuclear activity.  Indeed, not only the triggering but also the fueling of this activity can be studied in these objects. Since their orientation provides continuum emission on both sides of the center of activity one can search for either H\\kern0.1em{\\sc i}\\ or free-free absorption (Peck et al. \\cite{pec00b}).  The radio galaxy 4C\\,31.04 (B2\\,0116+31) was observed with VLBI by Wrobel \\& Simon (\\cite{wro86}) at 92 cm using 7 antennas.  Their image showed two components separated by about 70 mas. Cotton et al. (\\cite{cot95}) tentatively classified this source as a low redshift ($z=0.0592$) CSO. Giovannini et al. (2001, hereafter \\cite{gio01}) confirmed the CSO structure and identified the core with a weak flat spectrum central component.  On the basis of spectral line VLBI observations, Conway (1996 -- hereafter \\cite{con96}, \\cite{con99}) has presented evidence for a circumnuclear H\\kern0.1em{\\sc i}\\ disk, viewed close to edge-on, whose axis is that of the radio jet. The presence of a circumnuclear disk is also supported by HST observations (Perlman et al. \\cite{per01}), which reveal a disk-like dusty feature of somewhat larger scale but similar orientation to the one suggested by Conway.  The optical nucleus is also found to have cone-like features well aligned with the radio axis.  After a five year gap following the first observation at 5 GHz presented in \\cite{gio01}, we re-observed this source with the VLBA in order to look for a possible expansion. After describing the observations and data reduction in Sect.  2, we present our detailed comparison of the two epochs and information on spectral index in Sect.  3.  Some implications of the results are discussed in Sect.  4. Conclusions are summarized in Sect.  5.  We assume H$_0 = 65$ km s$^{-1}$ Mpc$^{-1}$ and q$_0 = 0.5$ throughout. At the distance of 4C\\,31.04 ($z=0.0592$) 1 mas corresponds to 1.2 pc and 1 mas/yr = 3.9 $c$. ",
        "conclusions": "From the comparison of the two VLBI images at 5 GHz, a small but significant expansion appears to be present in the Compact Symmetric Object 4C\\,31.04. From the derived expansion rate and the current size of the source we calculate an age of 550 years. A somehow larger value of the age (3000 -- 5000 yrs) is derived from spectral index analysis. Despite a significant difference, both results agree that the source is very young, as already demonstrated for other CSOs in the literature. However, the relatively high hot spots velocity and its low power make of 4C\\,31.04 a peculiar case among CSOs.  The discrepancy and uncertainty in the age estimates, the uncommon presence of a deep `hole' in the East lobe, the detection and the variability of the compact core suggest that it may be worth observing this source again in the future. New VLBI observations will be of interest to look for free--free absorption at low frequency, to confirm the expansion, and to look for spectral aging at high frequency."
    },
    "0212/astro-ph0212004_arXiv.txt": {
        "abstract": "Recent theoretical work has made it plausible for neutron stars (NSs) to lose angular momentum via gravitational radiation on long timescales ($\\sim10^6$ yr) while actively accreting. The gravitational waves (GWs) can either be emitted via the excitation of r-modes or from a deformed crust. GW emission can thus intervene to slow-down or halt the otherwise relentless spin-up from accretion. Prior to this theoretical work (and the measurements of NS rotation rates in LMXBs, see Chakrabarty's contribution) the community was rather confident that an accreting NS would be spun-up to rotation rates near breakup, motivating searches for sub-millisecond objects. After only briefly describing the physics of the GW processes, I argue that the limiting spin frequency might be appreciably lower than the breakup frequency. Millisecond radio pulsar observers would likely discover the impact of GW emission as a dropoff in the number of pulsars beyond 600 Hz, and I show here that the millisecond pulsar inventory in 47 Tuc might already exhibit such a cutoff. These theoretical ideas will be tested by GW searches with ground-based interferometers, such as the advanced LIGO instrument proposed for operation by 2008. ",
        "introduction": "It has been over twenty years since the discovery of the first millisecond radio pulsar (MSP) by Backer et al. (1982). Since then, large numbers of MSPs have been found in globular clusters and in the field. The theoretical expectation is that these NSs are ``recycled'' and spun-up by prolonged accretion (see reviews by Phinney \\& Kulkarni 1994; Bhattacharya 1995).  If accreting material arrives with the specific angular momentum of a particle orbiting at the NS radius ($R=10 R_6 {\\rm km}$), then it takes $\\approx 10^8 \\ {\\rm yr}$ of accretion at a rate $\\dot M\\approx 10^{-9} M_\\odot \\ {\\rm yr^{-1}}$ for a $M=1.4M_{1.4} M_\\odot$ NS to reach $\\nu_s=500$ Hz from an initially low frequency (see Cook et al. 1994 for a thorough set of numbers).  Many binary scenarios (e.g. Burderi et al. 1999) easily provide such an $\\dot M$ and transfer enough material to spin-up the NS to submillisecond periods. Hence, {\\it if there is no limiting physics}, rapid rotation is expected and there should be no ``cutoff'' in the population other than that from the transferred mass.  The piece of limiting physics typically discussed is magnetic accretion (e.g. Ghosh \\& Lamb 1979).  For a given $\\dot M$ and magnetic dipole strength, $\\mu$, the NS has a magnetospheric radius inside of which the flow is fixed by the magnetic topology. However, as the NS is spun-up, material rotating in the disk at the magnetospheric radius eventually co-rotates with the NS, such that further accretion will not likely spin it up.  Namely, an equilibrium spin period is reached that matches the Kepler period at the magnetosphere \\begin{equation} P_{eq}=8\\ {\\rm ms}\\left({10^{-9} M_\\odot \\ {\\rm yr^{-1}}}\\over \\dot M\\right)^{3/7}\\left(\\mu\\over 10^{27} {\\rm G \\ cm^{3}}\\right)^{6/7}, \\end{equation} and further accretion is likely accomodated by alternating epochs of spin-up and spin-down. Such phenomena have been seen in the highly magnetic ($\\mu>10^{30} \\ {\\rm G \\ cm^{3}}$, $B>10^{12} \\ {\\rm G}$) accreting pulsars in high mass X-ray binaries (see Bildsten et al. 1997 for an overview). The $P_{eq}$'s for these NSs range from 1-1000 seconds for $\\dot M\\approx 10^{-10}-10^{-7} M_\\odot \\ {\\rm yr^{-1}}$, clearly implying a range of $\\mu$'s.  \\begin{table}[htb] \\begin{center} \\caption{Oscillations During Type I X-Ray Bursts} \\begin{tabular}{llllll} \\hline Object Name & $\\nu_B$ & Flux & $h_c$ &\\\\ & (Hz) & $10^{-9}$(cgs) & ($10^{-27}$) & \\\\ \\hline 4U 1916-053 & 270 & 0.5 & 1.0 \\\\ 4U 1702-429 & 330 & 1.0 & 1.2 \\\\ 4U 1728-34 & 363 & 2.8 & 2.0 \\\\ KS 1731-260 & 524& 0.2-2 &1.3 \\\\ Aql X-1 & 549 & 1.0 & 1.0 \\\\ MXB 1658-298 & 567 & 0.1& 0.3\\\\ 4U 1636-53 & 581 & 4.4 & 2.0\\\\ MXB 1743-29 & 589 & ? &?\\\\ SAXJ1750.8-2980 & 601 &? &? \\\\ 4U 1608-52 & 619 & $<1$ & $<1$ \\\\ \\hline \\end{tabular} \\end{center} \\end{table}  As discussed by Chakrabarty, the $\\nu_s$'s inferred from nearly coherent oscillations (given by $\\nu_B$ in Table 1) during Type I bursts finally revealed rapidly rotating NSs in LMXBs. In addition, there are three accreting millisecond pulsars: \\sax \\ ($\\nu_s=401 \\ {\\rm Hz}$; Wijnands \\& van der Klis 1998), \\xtegsfc \\ ($\\nu_s=435 \\ {\\rm Hz}$; Markwardt et al. 2002) and \\xtemit \\ ($\\nu_s=185 \\ {\\rm Hz}$; Galloway et al. 2002) which have $\\sim 10^{-11} M_\\odot \\ {\\rm yr^{-1}}$. All of the $\\nu_s$'s are at least a factor of two away from the breakup frequency. It is also striking that, despite $\\dot M$ contrasts of nearly a factor of 1000, the NS rotation rates are all within a factor of 2-3 of each other. White and Zhang (1997) (and others) have argued that this similarity arises because these NSs are magnetic and have reached $P_{eq}$. This has two requirements: (i) a tight relation between $\\mu$ and $\\dot M$ so that they all reach the same equilibrium, and (ii) a method of hiding the persistent pulse typically seen from a magnetic accretor.  The state of knowledge of the $\\mu$ values of these accreting NSs is quite limited (see Cumming's contribution), though we desire them to have adequate $\\mu$'s to become MSPs.  The current $\\nu_s$ measurements of these accreting NSs thus suggests a limiting frequency that is not easily accomodated by either magnetic limitations or lifetime arguments. This indication, combined with the theoretical insights summarized next, points to GW emission as a new piece of physics that can limit the spin-up. This would be ``bad news'' for those carrying out searches for previously hidden large populations of submillisecond MSPs, but ``good news'' for ground-based gravitational wave searches. ",
        "conclusions": "I have hopefully made the case that both theory and observation point to a new limiting mechanism to NS spin-up: GW emission.  If true, the impact on MSP searches is clear: {\\it many fewer rapidly rotating MSPs will be found than predicted in those scenarios that neglect this physics.} One way that such physics could reveal itself is to put in a ``cutoff'' in the MSP population at some particular frequency. The prime complication is attributing such a cutoff to the underlying population and not observational selection.  However, in 47 Tuc, the observational selection is calculable, and between 300 and 600 Hz, the MSP sensitivity drops by only 30 percent (see Figure 2 of Camilo et al. 2000). Such a small sensitivity drop is not adequate to explain the population decline (see solid line in Figure 2) to high frequencies in 47 Tuc. I also plot the $\\nu_s$ distribution for 13 accreting NS with a dashed line. This sample is not, to our knowledge, biased in frequency in any way. If this is the injection distribution that becomes like that in 47 Tuc, then the higher $\\nu_s$ systems must spin-down a factor of 1.5 in 5-8 Gyrs, easily accomodated with a surface field of $B\\approx (2-5)\\times 10^8 {\\rm G}$.  \\noindent"
    },
    "0212/astro-ph0212468_arXiv.txt": {
        "abstract": "A major goal of contemporary astrophysics is understanding the origin of the most massive galaxies in the universe, particularly nearby ellipticals and spirals.  Theoretical models of galaxy formation have existed for many decades, although low and high redshift observations are only beginning to put constraints on different ideas.  We briefly describe these observations and how they are revealing the methods by which galaxies form by contrasting and comparing fiducial rapid collapse and hierarchical formation model predictions.  The available data show that cluster ellipticals must have rapidly formed at $z > 2$, and that up to 50\\% of all massive galaxies at $z \\sim 2.5$ are involved in major mergers. While the former is consistent with the monolithic collapse picture, we argue that hierarchal formation is the only model that can reproduce all the available observations. ",
        "introduction": "Massive galaxies, typically those with stellar masses $> 10^{10}$ M$_{\\odot}$, are the best studied galaxies at all distances due to their high luminosities.  Because we can study these galaxies in detail, nearly all galaxy formation ideas, models, and scenarios are geared towards constraining the properties of these systems (Tinsley \\& Gunn 1976; Cole et al. 2000). A persistent unanswered question is however: how did these galaxies form?   There are currently no lack of theoretical answers, although observations are only just beginning to constrain different ideas.  It is very popular to divide major galaxy formation scenarios into two different classes: the so-called monolithic collapse (Tinsley \\& Gunn 1976) and hierarchical formation, usually in the milieu of the Cold Dark Matter (CDM) paradigm (Cole et al. 2000).   If we accept that galaxies are embedded in dark matter halos then  we have an additional, and perhaps more fundamental, problem of determining how dark halos assembled, and how galaxy formation occurred in concert with this.  Monolithic collapse models assume that early in the universe's history ($z > 2$) baryonic gas in galaxies collapsed to form stars within a very short period of time ($\\sim 100$ Myr) with star formation rates of 10$^{2}$ - 10$^{3}$ \\solm yr$^{-1}$,  creating massive galaxies that thereafter passively evolve in luminosity.  In hierarchical models, galaxies form through mergers of pre-existing smaller systems. With these two major ideas as guidelines, we will describe observational properties of the most massive galaxies in the nearby and distant universe and how these observations currently favor the hierarchical idea.  Currently we believe that giant elliptical galaxies in clusters formed their stellar mass early in the universe at $z > 2$ (e.g., Ellis et al. 1997) and most galaxy formation studies also try to answer when these objects formed. This review is focused on answering the other fundamental origin question of {\\it how} this formation occured.  Three basic pieces of evidence are discussed: 1. The fact that nearby giant ellipticals contain old ($> 8$ Gyr) stellar populations, 2. The well-defined scaling relations for ellipticals (i.e., fundamental plane) seen out to $z \\sim 1$, and 3. The structural appearances of high-$z$ galaxies.  While monolithic collapse models can account for the first two through a rapid early formation, the spatial structures of high-$z$ galaxies at $z \\sim 2.5$ can only be reconciled with the hierarchical model. \\vspace*{-0.5cm} ",
        "conclusions": ""
    },
    "0212/astro-ph0212142_arXiv.txt": {
        "abstract": "We present {\\it HST}/NICMOS [Fe\\,{\\sc ii}]$1.644\\,\\mu$m, Pa$\\alpha$ ($1.87\\,\\mu$m) and continuum images of the starburst galaxies M82 and NGC~253 at an unprecedented spatial resolution. In both galaxies we detect [Fe\\,{\\sc ii}] compact sources superimposed on a diffuse background in the disk of the galaxies together with a component above and below the plane of the galaxy. The radio and [Fe\\,{\\sc ii}] emissions perpendicular to the disk of M82 show a remarkable similarity to each other, suggesting that both emissions originate in shocks from supernova explosions. We find a spatial correspondence between bright {\\it compact} [Fe\\,{\\sc ii}] emitting regions and the location of radio supernova remnants (SNR) for approximately $30-50$\\% of radio SNRs in  M82 and NGC~253. This lack of a one-to-one correspondence, more than being indicative of a different origin for the radio and [Fe\\,{\\sc ii}] emission in starbursts,  suggests two populations of SNRs: an older population ($ \\lesssim 10^4\\,$yr) traced by the [Fe\\,{\\sc ii}] emission and a younger population (a few hundred years old) traced by the radio SNRs. We therefore conclude that the [Fe\\,{\\sc ii}] emission in starburst galaxies provides a good estimate of the supernova activity. Using our newly determined [Fe\\,{\\sc ii}] luminosities (corrected for extinction) of M82 and NGC~253, we reevaluate the calibration of the supernova rate in terms of the [Fe\\,{\\sc ii}] luminosity for starburst galaxies. ",
        "introduction": "Compact radio sources in nearby starburst galaxies often appear to be supernova remnants (SNR). SNR in our Galaxy and the Large Magellanic Cloud (LMC)  are sources of bright near-infrared [Fe\\,{\\sc ii}] line emission with luminosities ranging from $10^{26}\\,$W to $3\\times10^{29}\\,$W (for the [Fe\\,{\\sc ii}]$1.644\\,\\mu$m line, e.g., Oliva, Moorwood, \\& Danziger 1989; Burton \\& Spyromilio 1993). This behavior suggests a connection between supernova activity and the [Fe\\,{\\sc ii}] emission in starbursts. This is supported by near-infrared spectroscopy (Moorwood \\& Oliva 1988; Greenhouse et al. 1991; Vanzi \\& Rieke 1997), by high-resolution [Fe\\,{\\sc ii}] imaging studies (Forbes et al. 1993; Greenhouse et al. 1997; Alonso-Herrero et al. 2000), and by the tight correlation between the radio 6\\,cm and the [Fe\\,{\\sc ii}] emissions (Forbes \\& Ward 1993). Further determination of the relation of the infrared [Fe\\,{\\sc ii}] emission with the supernova activity in nearby starforming galaxies would be helpful both for the understanding of these galaxies and for using the near-infrared [Fe\\,{\\sc ii}] luminosity as estimator of the supernova rate in more distant galaxies where individual radio supernovae cannot be identified.  There are indications of some complexities in this relation. For example, the bright [Fe\\,{\\sc ii}] emission may not be spatially coincident with the radio SNR. Lumsden \\& Puxley (1995) conducted a spectroscopic survey for [Fe\\,{\\sc ii}] emission from radio SNRs in M33 and found that in some cases the [Fe\\,{\\sc ii}] emission does not peak at the position of the radio SNR, as well as that some radio SNR show no [Fe\\,{\\sc ii}] emission. Recently, Morel, Doyon, \\& St-Louis (2002) have carried out an [Fe\\,{\\sc ii}] imaging survey of optically selected SNR in M33 and found a similar result. Both works have concluded that the [Fe\\,{\\sc ii}] emission from SNR may be produced over a time scale of $\\sim$ 10$^4$ years. Greenhouse et al. (1997) found that none of the brightest sources of [Fe\\,{\\sc ii}] emission in M82 were coincident with the position of the radio SNRs.  M82 and NGC~253 are the nearest starburst galaxies (distances $d = 3.2\\,$Mpc and $d = 2.5\\,$Mpc respectively). Both galaxies show a large number of radio compact sources (e.g., $\\simeq 50$ in M82 Kronberg et al. 1985 and Huang et al. 1994; $\\simeq 60$ in NGC~253 Turner \\& Ho 1985 and Ulvestad \\& Antonucci 1991, 1997 and references therein). Their proximity makes these two galaxies ideal to study the relation between supernova activity and the [Fe\\,{\\sc ii}] emission in starburst galaxies. Greenhouse et al. (1997) presented a Fabry-Perot [Fe\\,{\\sc ii}]$1.644\\,\\mu$m map  for M82 with a spatial resolution of $\\simeq 1.3 \\arcsec$. Their main conclusions were that the compact [Fe\\,{\\sc ii}] sources contribute  only $\\simeq 14\\%$ of the total [Fe\\,{\\sc ii}] emission in M82 and that they trace a population of SNR older than the radio SNR. Their work is somewhat inconclusive because of the  limited angular resolution and also because strong extinction could hide [Fe\\,{\\sc ii}] emission from some of the radio SNRs. This paper presents {\\it HST}/NICMOS images of M82 and NGC~253. The high spatial resolution of the NICMOS  [Fe\\,{\\sc ii}]$1.644\\,\\mu$m images (${\\rm FWMH} = 0.35\\arcsec$; this corresponds to  5\\,pc and 4\\,pc for M82 and NGC~253, respectively) allows a detailed study of the location of the bright sources of [Fe\\,{\\sc ii}]  and of the r\\^ole of supernova remnants in this line emission. ",
        "conclusions": "We have presented {\\it HST}/NICMOS observations of the starburst galaxies M82 and NGC~253. The morphology of the [Fe\\,{\\sc ii}]$1.644\\,\\mu$m line emission in both galaxies displays two components: (1) a disk component with a number of superimposed compact sources  and (2) a diffuse component in the form of filaments above and below the disk of the galaxy. In M82 we find a remarkable similarity between the extended component, the fingers or chimneys  above and below the disk, of the radio and [Fe\\,{\\sc ii}] emissions.  In M82, we have detected in both Pa$\\alpha$ and [Fe\\,{\\sc ii}] a nuclear ring of star formation. This ring, with an approximate diameter of 400\\,pc, has also been observed in other near- and mid-infrared hydrogen recombination lines and in the [Ne\\,{\\sc ii}]$12.8\\,\\mu$m line. The peak of the $2.2\\,\\mu$m continuum brightness of M82 lies very close to (within 1\\arcsec) the dynamical center, and corresponds to the center of the ring of star formation. In NGC~253, bright Pa$\\alpha$ emission stems from the near-infrared continuum peak and extends in a plume-like morphology along the disk of the galaxy. There is also a diffuse Pa$\\alpha$ component above and below the plane of galaxy, however, much fainter than that in M82.  Prior to calibrating the surpernova rate in terms of the [Fe\\,{\\sc ii}] emission, we have derived the extinction to the central regions of NGC~253 and M82. To assess the extinction correction, we have compared  the number of ionizing photons infered from the extinction-corrected Pa$\\alpha$ maps with other estimates based on longer wavelength indicators, and found a good agreement.   Since the [Fe\\,{\\sc ii}] emission in starburst galaxies is thought to be associated with supernova activity, we have cross-correlated the positions (with a spatial resolution of a few tenths of arcseconds) of the radio SNRs and the identified compact [Fe\\,{\\sc ii}] sources in M82 and NGC~253. We have found that approximately $30-50$\\% of the radio SNRs have an [Fe\\,{\\sc ii}] counterpart. However, we cannot associate an [Fe\\,{\\sc ii}] source to each radio SNR or vice versa. This is well understood in terms of the different ages traced by the radio SNRs (a few hundred years) and the [Fe\\,{\\sc ii}] stage of SNRs ($ \\lesssim 10^4\\,$yr). The compact [Fe\\,{\\sc ii}] emitting sources in M82 and NGC~253 only contribute to some $20-30\\%$  of the total [Fe\\,{\\sc ii}]$1.644\\,\\mu$m emission.  However, much of the remaining emission in the disk of the two galaxies is most likely produced by SNRs that expanded and merged into a general ISM a few $10^4\\,$yr ago. Photometry on the [Fe\\,{\\sc ii}]$1.644\\,\\mu$m emission arising from the disk of M82 shows that as much as $70\\%$ of the total [Fe\\,{\\sc ii}]$1.644\\,\\mu$m emission (corrected for extinction) is associated with SNRs. The remaining diffuse [Fe\\,{\\sc ii}] emission perpendicular to the disk of M82 seems to be related to  the superwind (via expulsion of material from the central starburst), as argued from the radio emission morphology (Wills et al. 1999).   All these arguments support the supernova origin for the observed near-infrared [Fe\\,{\\sc ii}] line emission in starburst galaxies. Thus we have recalibrated the supernova rate in terms of the [Fe\\,{\\sc ii}] luminosity relation for starburst galaxies using our newly measured (corrected for extinction) [Fe\\,{\\sc ii}]$1.644\\,\\mu$m luminosities of M82 and NGC~253."
    },
    "0212/astro-ph0212412_arXiv.txt": {
        "abstract": "We present X-ray imaging, timing, and phase resolved spectroscopy of the anomalous X-ray pulsar \\fouru\\ using the \\cxo\\ (CXO). The spectrum is well described by a power law plus blackbody model with $\\Gamma = 3.35(2)$, \\kt$=0.458(3)$~keV, and \\nh$=0.91(2)~\\times~10^{22}~{\\rm cm}^{-2}$; we find no significant evidence for spectral features ($0.5-7.0$ keV). Time resolved X-ray spectroscopy shows evidence for evolution in phase in either $\\Gamma$, or \\kt\\, or some combination thereof as a function of pulse phase. We derive a precise X-ray position for the source and determine its spin period, $P=8.68866(30)$~s. We have detected emission beyond $4\\arcsec$ from the central source and extending beyond $100\\arcsec$, likely due to dust scattering in the interstellar medium. ",
        "introduction": "\\label{sec:intro} \\fouru\\ was discovered in 1978 with {\\em Uhuru} \\citep{Forman78}, but only in 1984 were 8.7-s pulsations detected in its persistent X-ray flux with {\\em EXOSAT} \\citep{Israel94}. Pulse frequencies were sparsely measured until the last few years.  The frequency values are consistent with a constant spin-down rate of $\\dot \\nu = (-3.0\\pm 0.1) \\times 10^{-14}$ Hz s$^{-1}$, slightly larger than the $\\dot\\nu = (-2.598 \\pm 0.002) \\times 10^{-14}$ Hz s$^{-1}$ recently measured with {\\em RXTE} phase connected data \\citep{Gavriil01}. Despite detailed searches \\citep{Israel94,Wilson99}, no evidence for orbital Doppler shifts has been found in the pulse frequency measurements. Upper limits of $a_{\\rm X} \\sin i \\lesssim 0.26$ lt-s for orbital periods in the range 70 s to 2.5 days \\citep{Wilson99} place stringent limits on allowed orbital inclinations and companion masses for normal or helium main sequence companions.  {\\em ASCA} observations in 1994 \\citep{White96} and in 1998 \\citep{Paul00}, both determined the source energy spectrum as an absorbed power law (photon index $\\sim 3.7$) with a blackbody ($kT \\sim 0.4$ keV) component. The {\\em ASCA} GIS data also revealed very stable (in time) pulse profiles and intensities. The pulse profiles, however, evolved with energy from double-peaked (0.5-4.0 keV), to single peaked (4.0-8.0 keV) \\citep{White96}. Later spectral observations with {\\em BeppoSAX} in 1997-1998 \\citep{Israel99} were consistent with the {\\em ASCA} spectral results, however it did not show significant pulse phase dependence.  Using RXTE observations, \\cite{Gavriil01} reported pulse profile evolution with energy, evolving from double-peaked (2.0$-$4.0 keV) to single peaked (6.0$-$8.0 keV). \\cite{Juett02} used the High Energy Transmission Grating Spectrometer (HETG) on the CXO to observe the source in 2001.  They find no evidence for emission or absorption lines in the spectrum which is well fitted by an absorbed powerlaw$+$blackbody with $\\Gamma=3.3(4)$ and $kT=0.418(13)$ keV. Their analysis used only phase averaged spectra and provided stringent constraints on spectral features.  Recent optical observations \\citep{Hulleman00} have revealed an object with peculiar optical colors inside the {\\em Einstein} error circle \\citep{White87}. The optical observations imply an extinction of $2.7 \\lesssim A_{V} \\lesssim 5.0$ and a distance of $d \\gtrsim 2.7$ kpc. The object is too faint to be a large accretion disk. Further observations of the optical counterpart revealed pulsations with a 30\\% pulse fraction \\citep{Kern01}. The large pulse fraction also rules out X-ray reprocessing (in a disk) as the source of the optical pulsations. A deep radio search showed no evidence for a SNR, extended structure, or even a point source at the position of \\fouru\\ \\citep{Gaensler01}.  We present here an analysis and discussion of our \\chan\\ observation of the source. In \\S \\ref{sec:obsresults} we describe the x-ray observations, and present spatial, timing and spectral results of \\fouru.  We also include spatially resolved spectroscopy of the extended emission around \\fouru, presumably due to dust scattering.  A precise x-ray position is derived and compared with the optical identification from \\cite{Hulleman00}. ~ ",
        "conclusions": "\\label{sec:discuss}  We have observed the anomalous X-ray pulsar, \\fouru\\ with \\chan\\ and performed detailed phase resolved spectroscopy of the source. We find that there is significant spectral evolution during the overall pulse.  The phase averaged spectrum can be best fit with a two component model (PL+BB), consistent with earlier results (\\cite{Juett02,White96,Israel99,Paul00}). We confirm the source location of \\cite{Juett02}, which also coincides with the optical counterpart of \\cite{Hulleman00}. We measure a pulse period consistent with the measurements of \\cite{Gavriil01} and confirm limits to the amplitude of possible binary motion in a subset of period ranges previously investigated by \\citep{Wilson99}.  The superb \\chan\\ spatial resolution allows us to determine that there is emission extending up to $100\\arcsec$ from the pulsar, likely due to dust scattering. We are unable to determine evidence for excess emission within a few arcseconds of the source (i.e., evidence for a plerion), due to contamination from the X-ray bright AXP and the associated dust scattering halo."
    },
    "0212/astro-ph0212138_arXiv.txt": {
        "abstract": "Analyses of stellar spectra often begin with the determination of a number of  parameters that define a model atmosphere. This work presents a prototype  for an automated spectral classification system that uses a 150 \\AA-wide region around H${\\beta}$, and applies to stars of spectral types A to K with normal (scaled solar) chemical composition. The new tool exploits synthetic spectra based on plane-parallel flux-constant model atmospheres. The input data are high signal-to-noise spectra with a resolution greater than about 1 \\AA. The output parameters are forced to agree with an external scale of effective temperatures, based on the Infrared Flux Method. The system is fast -- a spectrum is classified in a few seconds-- and well-suited for implementation on a web server. We estimate upper limits to the $1\\sigma$ random error in the retrieved effective temperatures, surface gravities, and metallicities as  100 K, 0.3 dex, and 0.1 dex, respectively. ",
        "introduction": "Mass, radius, and luminosity  are some of the most interesting properties of a star. Unfortunately, it is non-linear combinations of them that produce quasi-linear changes on a stellar spectrum. Stellar fluxes are commonly interpreted in terms of atmospheric temperature, pressure, and chemical composition. In the context of classical flux-constant model atmospheres, these fields are simply specified with three scalars: the effective temperature, the surface gravity, and a solar-scaled metallicity. Nevertheless, extracting the three from an observed spectrum is rarely trivial.  A  spectroscopic classification system has been developed independently of the physical parameters. The MK system (Morgan, Keenan \\& Kellman 1943) lays out a series of rules to assign spectral classes from medium-low resolution spectra. This method has the advantage of providing a standard reference independent of models. However, it is somewhat artificial, in the sense that the defined spectral classes are obviously  correlated with the relevant atmospheric parameters. Moreover, the classical MK system does not provide for  metal-poor stars. One may note, as an example, that the metal-poor giant HD 122563, which has an effective temperature around 4600 K, has been often classified as a late-F or early-G type star. On the other hand, classification methods based on  physical parameters are more natural, but model-dependent.  Both the MK and the {\\it physical} classification systems have their own advantages,  and this may be the reason why they still coexist. But, the fact  that it is the physical parameters  what is ultimately demanded for astronomical applications is shifting most of the recent research toward the direct extraction of those quantities. As recognized by many specialists, repeatability and high speed in spectroscopic stellar classification can only be achieved by using automatic methods. A recent discussion  of the most used methods can be found in Bailer-Jones (2001).   Derived stellar parameters  may differ when determined from  different wavelength ranges or spectral features. Among other reasons, this may be caused by using models that are too simplistic. As the parameters will be derived from their expected effect on the spectra, inaccurate predictions, or neglect of other relevant parameters, will bias the results. Details in the implementation of a classification procedure are also a reason to worry. The wide range of effective temperatures that have been assigned to the metal-poor subgiant HD 140283 in the recent literature serves as testimony of the respectable uncertainties still involved in the scale of effective temperatures (see, e.g., the discussion in Snider et al. 2001).  Discrepancies produced by interpreting stellar spectra with model atmospheres that are too simple will decrease as progress in theory takes place. For main-sequence stars, remarkable advances are happening as  efforts focus on relaxing the assumptions of Local Thermodynamic Equilibrium (LTE) and hydrostatic equilibrium. Systematic differences that arise in the implementation of different methods for spectroscopic classification can be controlled by establishing standards. A reference implementation would be far more powerful than a  set of standard stars. Modern information technologies have paved the way for implementing a public automatic classification system accessible over the Internet. Such open system, if reliable and fast,  could serve as a standard reference.  \\begin{figure*} \\centering \\includegraphics[width=8cm,angle=90]{f1.ps} \\protect\\caption[ ]{ Observed fluxes for Procyon, the Sun and Arcturus in the selected window. The original very high dispersion observations have been convolved with a Gaussian profile with a FWHM of 1.0 \\AA. The observations in the McDonald atlas of Procyon (Allende Prieto et al. 2002) have been continuum normalized by P. S. Barklem (see Barklem et al. 2002 for more details). A missing interval in the spectrum of Arcturus is severely affected by telluric lines. \\label{f1}} \\end{figure*}  This paper discusses the first steps toward the implementation of a prototype for an open classification system. This work will only deal with a section of the HR diagram; stars with spectral types A to K and scaled solar metal abundances. Section 2 describes the selected wavelength  band,  \\S 3 the working parameters space, and \\S4 provides details of the implementation. Section 5 is devoted to connecting the parameters derived from spectroscopy to a widely accepted scale of effective temperatures based on the Infrared Flux Method, and checking the performance of the method. Section 6 concludes with a summary of present results and ideas for subsequent work. ",
        "conclusions": "We have implemented a spectroscopic classification algorithm that provides estimates for $T_{\\rm eff}$, $\\log g$, [Fe/H], and $\\xi$ based on the observations with a resolving power $R \\ge 5000$ in the spectral range 4810--4960 \\AA. The classification system is based on synthetic spectra calculated with classical flux-constant model atmospheres, and it is anchored to the photometric calibrations of effective temperature derived by Alonso and collaborators (Alonso et al. 1996, 1999) based on the Infrared Flux Method. By using the Elodie spectroscopic library (Prugniel \\& Soubiran 2001) and the gravities determined by Allende Prieto \\& Lambert (1999) for nearby Hipparcos stars, we derive upper limits to the uncertainties in the retrieved $T_{\\rm eff}$, $\\log g$, and [Fe/H], as 100 K, 0.3 dex, and 0.1 dex, respectively.  Our classification algorithm  can be used on spectral types A to K, for main-sequence and evolved stars with gravities as low as $\\log =2$, and all metallicities. The system is able to classify a stellar spectrum in only 3 seconds on a modern workstation. Work is in progress to improve the accuracy of the synthetic spectra by using  a more realistic line absorption profile for H${\\beta}$. We also plan on varying the parameters that affect the performance of the employed GA, and testing different optimization algorithms. An important question to answer is how the performance of our system improves or degrades with spectral resolution and signal-to-noise ratio. Different spectral ranges should be explored. Spectral bands with lines of metals whose abundances do not scale well with iron should be avoided, unless more parameters are included in the search, which is certainly feasible.  Different classification algorithms for stellar spectra have been tested in the literature. In particular, artificial neural networks hold the promise for the highest speeds, which may be critical for problems involving a large number of free parameters (see, e.g. Bailer-Jones 2001; Snider et al 2001). When the number of parameters is limited, like in the spectroscopic classification considered here, GAs have the advantage of not requiring training. This, in turn, allows us to explore different strategies, such as the selection of the spectral range or the resolving power, very quickly, while keeping the search global.  Our final goal is to provide a web interface and make this or a similar system publicly available. Future  work will target hotter spectral types, using calculations based upon non-LTE model atmospheres. Extension to the bottom of the main-sequence and beyond could follow, although theoretical modelling of cool atmospheres  has not yet reached the same level of maturity as for warmer stars. The adopted strategy can easily accommodate future improvements in model atmospheres and spectral synthesis."
    },
    "0212/astro-ph0212248_arXiv.txt": {
        "abstract": "\\vspace{\\baselineskip} Simple curvaton models can generate a mixture of of correlated primordial adiabatic and isocurvature perturbations. The baryon and cold dark matter isocurvature modes differ only by an observationally null mode in which the two perturbations almost exactly compensate, and therefore have proportional effects at linear order. We discuss the CMB anisotropy in general mixed models, and give a simple approximate analytic result for the large scale CMB anisotropy. Working numerically we use the latest WMAP observations and a variety of other data to constrain the curvaton model. We find that models with an isocurvature contribution are not favored relative to simple purely adiabatic models. However a significant primordial totally correlated baryon isocurvature perturbation is not ruled out. Certain classes of curvaton model are thereby ruled out, other classes predict enough non-Gaussianity to be detectable by the Planck satellite.  In the appendices we review the relevant equations in the covariant formulation and give series solutions for the radiation dominated era. ",
        "introduction": "Recent detailed measurements of the acoustic peaks in CMB anisotropy power spectrum by the WMAP satellite~\\cite{Hinshaw03,Peiris03} are consistent with the standard model of a predominantly {\\em adiabatic}, approximately scale invariant primordial power spectrum in a spatially flat Universe. Frequently it is assumed the initial power spectrum is entirely adiabatic, though there is still no compelling justification for this assumption. Although adiabatic perturbations are predicted from single field models of inflation \\cite{LL}, if one allows the possibility of multiple fields in the early Universe then there is also the possibility of \\emph{isocurvature} perturbations  (also known as \\emph{entropy} perturbations) \\cite{Kofman87,Polarski94,Garcia-Bellido96,Linde97,Sasaki98,Langlois99,Gordon00,Hwang00,Starobinsky01,Bartolo01,GrootNibbelink01,Wands02,Tsujikawa02}. In particular, the recently proposed {\\em curvaton} model uses a second scalar field (the `curvaton')  to form the perturbations \\cite{Mollerach90,Lyth01,Lyth02,Moroi01,Moroi02, Murayam02}. The motivation for this is it makes it easier for otherwise satisfactory particle physics models of inflation to produce the correct primordial spectrum of perturbations \\cite{Dimopoulos01}. Various candidates for the curvaton have been proposed \\cite{Postma02,Enqvist02,Bastero-Gil02,Dimopoulos03a,Dimopoulos03b}. A curvaton mechanism has also been considered in the pre big bang scenario \\cite{Copeland97,Lidsey00,Enqvist01,Sloth02,Bozza02} where it can be used to produce an almost scale invariant spectrum.  The curvaton scenario also has the feature of being able to generate isocurvature perturbations of a similar magnitude to the adiabatic perturbation without fine tuning, and therefore is open to observational test.  Early studies of non-adiabatic perturbations, either considered purely isocurvature cold dark matter perturbations \\cite{Efstathiou86} or mixtures of adiabatic and uncorrelated cold dark matter isocurvature perturbations \\cite{Stompor96,Pierpaoli99,Kawasaki2001,Enqvist2000}. However, as first realized by Langlois \\cite{Langlois99}, the adiabatic and isocurvature components can be correlated and this correlation may have interesting observational consequences \\cite{Langlois00}. In Ref.~\\cite{Bucher99} they identified four regular isocurvature modes, which in general can have arbitrary correlations with each other and with the adiabatic mode. Such general models have many degeneracies and are badly constrained by pre-WMAP data \\cite{Trotta01,Trotta02}. Detailed CMB polarization data is expected to help with this \\cite{Bucher01}. In Ref.~\\cite{Peiris03} (following Ref.~\\cite{Amendola01} with pre-WMAP data) they considered a CDM isocurvature mode with an arbitrary correlation to an adiabatic mode and found that though not favored by the data, a significant isocurvature contribution was still permitted. Constraints on a specific model that doesn't produce isocurvature modes were given in Ref.~\\cite{Bartolo02}.  Here we start in Section~\\ref{CMB} by making some general remarks about mixed isocurvature models, and discuss the corresponding CMB power spectra predictions. Then in Section~\\ref{curvaton} we discuss current observational constraints on totally correlated (or anti-correlated) adiabatic and isocurvature perturbations, as predicted by the curvaton model. Various scenarios within the curvaton model predict specific ratios of adiabatic and isocurvature perturbations, and can be tested directly. In general we find constraints on when the curvaton decayed.  We use the CMB temperature and temperature-polarization cross-correlation anisotropy power spectra from the WMAP\\footnote{\\url{http://lambda.gsfc.nasa.gov/}}~\\cite{Verde03,Hinshaw03,Kogut03} observations, as well as seven almost independent temperature band powers from ACBAR\\footnote{\\url{http://cosmologist.info/ACBAR}}~\\cite{Kuo02} on smaller scales. In addition we use data from the 2dF galaxy redshift survey~\\cite{Percival02}, HST Key Project~\\cite{Freedman01}, and nucleosynthesis~\\cite{Burles01} using a slightly modified version of the CosmoMC\\footnote{\\url{http://cosmologist.info/cosmomc}} Markov-Chain Monte Carlo program, as described in Ref.~\\cite{cosmomc}.   For simplicity we assume a flat universe with a cosmological constant, uninteracting cold dark matter, and massless neutrinos evolving according to general relativity. ",
        "conclusions": "\\label{conclusions}  The curvaton model provides a simple scenario that can give rise to correlated adiabatic and isocurvature modes of similar size. The current data do not favor a large isocurvature contribution, but a significant amplitude is still allowed.  We point out that the CDM and baryon isocurvature modes differ only by the addition of an observationally null mode in which the two perturbations compensate. The CDM isocurvature mode can therefore be treated as a scaled baryon isocurvature mode. A simple analytical approximation for the effect of mixtures of large scale isocurvature and adiabatic perturbations on the CMB temperature anisotropy was given. Numerically, we found that the data was consistent at the  two sigma level with the presence of an effective correlated  baryon isocurvature perturbation of about 50\\% the magnitude of the adiabatic perturbation. The individual baryon and CDM isocurvature modes can be even larger if they compensate each other. Models in which either the baryon number or CDM was created before the curvaton dominated the energy density are ruled out unless counter-balanced by the other species being created by the curvaton decay. The levels of non-Gaussianity expected for the various scenarios were evaluated and in the case of the CDM being created before the curvaton decayed and the baryon number by the curvaton decay, could be high enough to detectable by the Planck satellite."
    },
    "0212/astro-ph0212554_arXiv.txt": {
        "abstract": "The optical afterglow of gamma-ray burst 020813 was observed for 3 hours with the LRIS spectropolarimeter at the Keck-I telescope, beginning 4.7 hours after the burst was detected by \\hete.  The spectrum reveals numerous metal absorption lines that we identify with two systems at $z=1.223$ and $z=1.255$.  We also detect an [\\ion{O}{2}] $\\lambda3727$ emission line at $z=1.255$ and we identify this galaxy as the likely host of the GRB.  After a correction for Galactic interstellar polarization, the optical afterglow has a linear polarization of 1.8--2.4\\% during 4.7--7.9 hours after the burst.  A measurement of $p = 0.80\\% \\pm 0.16\\%$ on the following night by Covino \\etal\\ demonstrates significant polarization variability over the next 14 hours. The lack of strong variability in the position angle of linear polarization indicates that the magnetic field in the jet is likely to be globally ordered rather than composed of a number of randomly oriented cells.  Within the framework of afterglow models with collimated flows, the relatively low observed polarization suggests that the magnetic field components perpendicular and parallel to the shock front are only different by about 20\\%. ",
        "introduction": "It is now generally assumed that the optical afterglows of gamma-ray bursts (GRBs) arise from synchrotron emission in the post-shocked gas of a relativistic blast wave \\citep[see, e.g.,][]{pir00, mes02}. There is credible evidence that GRBs are strongly collimated or beamed explosions \\citep[e.g.,][]{rho99,sph99,fra01}.  Synchrotron radiation is intrinsically polarized, and the synchrotron emission from a jet will have a net polarization if its axis does not coincide with the line of sight.  A natural consequence of models for beamed afterglows is time-variable polarization, which will depend on the viewing geometry as well as the degree of orientation of the magnetic field \\citep[e.g.,][]{ml99, gru99}, up to a maximum polarization of $\\sim20\\%$ \\citep{gl99,sari99}.  Optical imaging polarimetry has been performed for a handful of GRB afterglows to search for this signature.  The results have been either upper limits \\citep{hjo99, cov02a} or detections at $p \\approx 1-3\\%$ \\citep{cov99, wij99, rol00, cov02b}, except for an observation of GRB 020405 by \\citet{ber02} that detected $9.9\\% \\pm 1.3\\%$ polarization 1.3 days after the burst.  GRB 020813, a typical long-duration burst with a series of pulses, was detected by \\hete\\ at 2:44 UT on 2002 August 13 \\citep{vil02}, and an associated optical transient (OT) was promptly identified \\citep{fbp02}.  The rapid localization allowed us to begin spectropolarimetric observations less than 5 hours after the burst was initially detected.  These are the first spectropolarimetric observations obtained for a GRB afterglow.  Preliminary descriptions of this dataset were given by \\citet{pri02} and \\citet{bar02}. ",
        "conclusions": "Our observations, combined with those of \\citet{cov02c}, clearly show variability in the degree of polarization, with a nearly constant position angle. This already places strong constraints on the polarization models.  Random effects, such as emission from a small number of randomly oriented magnetic field cells \\citep{gw99} or amplification of the emission from one such cell by microlensing \\citep{lp98}, would predict that $p$ and $\\theta$ should change on the same timescale.  In principle, the jet geometry, together with an assumption of some level of orientation of the amplified magnetic field with respect to the shock, could predict the degree of polarization as function of time.  Models based on this scenario \\citep{gl99,sari99} show that the maximum degree of polarization occurs around the time of the jet break (\\tjet), when the opening angle of the jet is equal to the inverse Lorentz factor.  However, at that time the emission is observable from a significant fraction of the jet. This introduces order-unity uncertainties in the polarization predictions since the energy distribution profile as function of distance from the axis of the jet is unknown. In addition, a detailed description of the hydrodynamic evolution is not yet available; only the asymptotic scalings, well before and well after \\tjet, are understood.  Thus, it is not possible to make detailed, quantitative comparisons of polarization models with observations.  Despite these uncertainties, we can make general comparisons of the observations with the model proposed by \\citet{sari99}, which predicts sharp features in the polarization light curves, although more realistic assumptions on the hydrodynamics of jet spreading, and on the energy distribution within the initial jet, are expected to produce smoother features than those predicted by this model.  This model has three parameters. The line-of-sight offset to the jet determines the shape of the curve; \\tjet\\ sets the overall timescale; and an overall scale factor for $p$ is set by a function of the ratio of the magnetic field components parallel to and perpendicular to the shock front, which is assumed to be constant.  We denote the offset in units of the opening angle of the jet by $0<\\phi<1$.  (For $\\phi > 1$ the jet is not visible initially, corresponding to an ``orphan'' afterglow.)  Small offsets give rise to a single polarization peak, while large offsets yield three peaks, where the middle one has position angle perpendicular to the other two.  More realistic jet models may smooth the curve and eliminate this middle peak.  We note that the model of \\citet{gl99} predicts only two peaks since their calculations did not include the spreading of the jet.  Figure \\ref{model} demonstrates that the jet model parameters can be chosen so as to find at least a qualitative agreement with the data. In the figure, our measurements for the 5800--6800 \\AA\\ polarization, and the $V$-band measurement from \\citet{cov02c}, are compared with two sample models.  Initial analysis of the optical light curve suggested $\\tjet \\approx 3.5-5$ hours \\citep{bfh02, gh02}, although subsequent analysis (to be presented elsewhere) indicates that the light curves may be consistent with a later value of $\\tjet \\approx 10-11$ hours.  Given this uncertainty in deriving \\tjet\\ from the photometric data, our approach is to determine whether there are any values of \\tjet\\ that produce polarization light curves consistent with the observations.  The first model has a small offset $\\phi=0.3$ and predicts one peak of polarization. (Models with $0 < \\phi \\lesssim0.35$ would give similar results.)  This requires \\tjet\\ to occur during the Keck observations, at $\\sim 6$ hours.  However, this model somewhat underpredicts the decay of polarization between our measurements and those of Covino \\etal\\ A better fit is obtained for a medium-offset model with $\\phi=0.55$ and $\\tjet \\approx 16$ hours, where the polarization drops to zero, or perhaps even rotates the position angle by 90\\arcdeg, between the Keck and VLT observations. Under both interpretations, the overall level of polarization is less than the maximal possible level of $\\sim20\\%$.  This implies that the magnetic field perpendicular to the shock differs only by $15-20$\\% from that parallel to the shock front.  Models with other values of $\\phi$ have difficulty matching both the near-constancy of $p$ during the Keck observations (for $\\lambda>5800$ \\AA) and the large change in $p$ detected by \\citet{cov02c}.  The polarization behavior of GRB 020813 appears rather similar to that of GRB 990712, which varied between $p=3.0\\%$ at 0.4 days after the burst and 1.2\\% at $t=0.8$ days, at nearly constant position angle \\citep{rol00}.  We note that our models with $\\phi=0.3$ or $0.55$ could also be nearly consistent with the GRB 990712 data within its uncertainties.  Still, as noted by Rol \\etal, the lack of clear evidence for a jet break for that event casts some doubt on this interpretation.  Finally, we mention that the hints of wavelength dependence to the polarization of GRB 020813, which appear to be time-variable, can not be accommodated by these models.  A spectral break passing through the optical band at the time of our three measurements could produce such a signature, but there is no evidence for such a break.  To date, polarimetric observations of GRB afterglows have had rather sparse temporal coverage.  Further progress will require continuous polarization monitoring of bright afterglows, initiated as early as possible and extending well beyond \\tjet.  Also, our data highlight the need for color information in future polarization measurements. Obtaining such observations would pose some logistical difficulties, but the potential for improved understanding of the physical structure of the synchrotron jets from GRBs would make this effort extremely worthwhile."
    },
    "0212/astro-ph0212391_arXiv.txt": {
        "abstract": "We present a detailed study of the relation between circumnuclear dust morphology, host galaxy properties, and nuclear activity in nearby galaxies. We use our sample of 123 nearby galaxies with visible--near-infrared colormaps from the {\\it Hubble Space Telescope} to create well-matched, ``paired'' samples of 28 active and 28 inactive galaxies, as well as 19 barred and 19 unbarred galaxies, that have the same host galaxy properties. Comparison of the barred and unbarred galaxies shows that grand design nuclear dust spirals are only found in galaxies with a large-scale bar. These nuclear dust spirals, which are present in approximately a third of all barred galaxies, also appear to be connected to the dust lanes along the leading edges of the large-scale bars. Grand design nuclear spirals are more common than inner rings, which are present in only a small minority of the barred galaxies. Tightly wound nuclear dust spirals, in contrast, show a strong tendency to avoid galaxies with large-scale bars. Comparison of the AGN and inactive samples shows that nuclear dust spirals, which may trace shocks and angular momentum dissipation in the ISM, occur with comparable frequency in both active and inactive galaxies. The only difference between the active and inactive galaxies is that several inactive galaxies appear to completely lack dust structure in their circumnuclear region, while none of the AGN lack this structure. The comparable frequency of nuclear spirals in active and inactive galaxies, combined with previous work that finds no significant differences in the frequency of bars or interactions between well-matched active and inactive galaxies, suggests that no {\\it universal} fueling mechanism for low-luminosity AGN operates at spatial scales greater than $\\sim 100$ pc radius from the galactic nuclei. The similarities of the circumnuclear environments of active and inactive galaxies suggests that the lifetime of nuclear activity is less than the characteristic inflow time from these spatial scales. An order-of-magnitude estimate of this inflow time is the dynamical timescale. This sets an upper limit of several million years to the lifetime of an individual episode of nuclear activity. ",
        "introduction": "Many observational programs over the past few years have led to the proposition that all galaxies with a substantial spheroid component contain supermassive black holes, irrespective of the presence or absence of nuclear activity \\citep[\\eg][]{richstone98,ferrarese00,gebhardt00}. Since black holes grow via the accretion of matter and this accretion leads to detectable nuclear activity, these results imply that all galaxies must go through an accretion phase, yet the mechanism which triggers nuclear activity in ``currently'' active galaxies remains unknown. In order to fuel active galactic nuclei (AGN), essentially all of the angular momentum must be removed from some fraction of the host galaxy's interstellar medium (ISM). Low-luminosity AGN, which dominate the local population, require accretion rates of $0.01 - 0.1 M_\\odot$ yr$^{-1}$, assuming typical radiative efficiencies.  Studies of AGN and inactive control samples have investigated the frequency of several mechanisms for angular momentum transport to determine their viability. Interactions between galaxies is one good candidate \\citep{toomre72} as theoretical simulations of mergers show significant accretion into the central regions of the merger remnant \\citep[\\eg][]{hernquist89,barnes91, mihos96}. Interactions may be responsible for triggering AGN activity in the more luminous quasars \\citep{hutchings97,bahcall97}, yet detailed studies of interacting pairs have not found a statistically significant excess of the lower-luminosity Seyfert galaxies in interacting systems \\citep{fuentes88}. Large-scale bars have also been proposed as a mechanism to fuel nuclear activity \\citep{simkin80,schwarz81}. The nonaxisymmetric potential due to a large-scale bar leads to the formation of a shock front along the bar's leading edges \\citep{prendergast83,athanassoula92} and material has been observed flowing into the central regions of several barred galaxies \\citep{quillen95,regan97,jogee99}. However, detailed near-infrared (NIR) studies of large samples of active and inactive galaxies have shown either no, or at most a marginal ($2.5\\sigma$), excess of large-scale bars in active samples \\citep{mulchaey97b,knapen00}. These studies of interacting and barred galaxies pushed the effective spatial resolution limit of ground-based observations for large samples of AGN, yet the typical spatial resolution of these investigations remain many hundreds of parsecs.  Several \\hst programs over the past few years have targeted the circumnuclear morphology of large active galaxy samples to search for signatures of AGN fueling \\citep[\\eg][]{malkan98,regan99,martini99}. One of the main goals of these programs was to investigate the fraction of Seyferts with nuclear bars (bars with semimajor axis lengths typically less than a kiloparsec), which could be comprised of gas or stars \\citep{shlosman89,pfenniger90} and cause the transport of matter from approximately a kiloparsec to tens of parsecs. However, these studies have found nuclear bars in only $\\sim 25$\\% of all Seyferts \\citep{regan99,martini01} and studies of Seyfert and control samples have found similar fractions of double bars in samples of active and inactive galaxies with large-scale bars \\citep{marquez00,laine02,erwin02}. The comparable fractions of nuclear bars in active and inactive galaxies, combined with the apparent absence of them in the majority of all active galaxies, suggests that some other mechanism is needed to fuel nuclear activity in many active galaxies.  One new candidate that arose from the \\hst studies is nuclear dust spirals \\citep{regan99,quillen99,martini99,pogge02}. Visible--NIR color maps of the majority of the active galaxies in these surveys showed nuclear spirals, with a wide range of coherence, that extend from approximately a kiloparsec down to tens of parsecs (the limiting spatial resolution of the nearest subsample). These nuclear spirals are distinct from the spiral arms in the main disks of these galaxies as they appear to have density contrasts of only a factor of a few above the ambient ISM and no associated star formation. Nuclear spirals are a promising fueling mechanism not only by virtue of their frequency, but also because they may mark the location of shock fronts or turbulence in the central, circumnuclear gaseous disks and therefore trace the sites of angular momentum dissipation. The possibility of shock-driven inflow, as traced by nuclear spiral structure, has been the subject of a number of recent theoretical studies \\citep{fukuda98,montenegro99,wada01a,maciejewski02,wada02}.  While most of the observational programs to date have targeted the circumnuclear region of active galaxies, nuclear dust spirals have also been observed in a small number of inactive galaxies with single-bandpass observations \\citep{phillips96,carollo98}. In \\hst Cycle~9 we began a program (SN8597, PI Regan) to obtain WFPC2 images of galaxies with prior NICMOS observations (from SN7330, PI Mulchaey and GO7867, PI Pogge) in order to quantify the frequency of nuclear spiral structure in inactive galaxies. We present the observations of our final sample of 123 galaxies, along with a description of the sample, survey design, and classification system for circumnuclear dust structure, in \\citet[][hereafter Paper~I]{martini02}. Our nuclear dust classification has six types, including four for nuclear dust spirals: grand design, tightly wound, loosely wound, and chaotic spirals. We placed galaxies with dust structures but without evidence for nuclear spirals in a fifth, ``chaotic'' class, and galaxies with no detected circumnuclear dust structure into a sixth, ``no structure'' class. The final dataset presented in Paper~I, in spite of the initial effort to create a well-match active and control sample, is relatively heterogeneous due to both the vagarious \\hst snapshot scheduling and our attempt to augment the sample with additional nearby galaxies of interest.  In the present paper we create well-matched subsamples of the full dataset presented in Paper~I in order to measure the relative frequency of nuclear dust spirals in active and inactive galaxies. This sample creation, described in the next section, draws unique pairs of active/inactive or barred/unbarred galaxies that otherwise have all of the same host galaxy properties and are at the same distance, while also maximizing the total size of the subsample drawn from the full dataset. In \\S\\S\\ref{sec:active} and \\ref{sec:all} we describe the trends for circumnuclear dust morphology for the active/inactive and barred/unbarred samples, respectively. The implications of these results for the fueling of nuclear activity are described in \\S\\ref{sec:discuss} and our results are summarized in \\S\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  The cold ISM in the circumnuclear region of nearly all spiral galaxies exhibits a wealth of structure. This structure often takes the form of nuclear dust spirals, which demonstrates that some differential rotation is present in the circumnuclear region of most spiral galaxies.  A comparison of a well-matched sample of 19 barred and 19 unbarred galaxies has confirmed our previous result (Paper~I) that grand design nuclear dust spirals are only found in galaxies with large-scale bars. Grand design spirals are present in approximately a third of all barred galaxies and these circumnuclear spirals connect to the relatively straight dust lanes associated with the large-scale bar. The dust lanes therefore form essentially contiguous structures from many kiloparsec scales to within tens of parsecs of galactic centers. These nuclear spirals are also more common than inner rings, which are found in only a minority of all barred galaxies. The fact that most barred galaxies do not have rings inside the bar suggests that bar-driven inflow does not commonly stall in a circumnuclear ring. Our comparison of barred and unbarred galaxies also shows that there is a marginally significant tendency for tightly wound nuclear dust spirals to be found in galaxies that lack large-scale bars. The low pitch angle of these spirals may be due to an increased time for winding or greater differential rotation.  To test for differences in the circumnuclear dust distribution in active and inactive galaxies, we have created a well-matched sample of 28 active and 28 inactive galaxies by identifying pairs with all of the same host galaxy properties. This analysis finds only a statistically insignificant excess of nuclear spirals in active galaxies over the inactive control sample and that nuclear spiral structure is present in the majority of both types. A perhaps more surprising result was that although all of the active galaxies have some discernible dust structure, seven of the 28 inactive galaxies have no circumnuclear dust structure.  These galaxies either have an extremely smooth dust distribution or no cold component to their ISM in the circumnuclear region. The comparable frequency of nuclear dust spirals in a well-matched sample of active and inactive galaxies, when combined with similar results for the frequency of bars or interactions from previous work, suggests that there is no universal fueling mechanism for low-luminosity AGN at scales greater than 100 pc from galactic nuclei.  We have shown that the circumnuclear dust morphologies of most active and inactive galaxies are indistinguishable. All of these galaxies are assumed to contain supermassive black holes and therefore the main difference between them is whether or not that central black hole is currently fueled. If the structures responsible for fueling nuclear activity are not present on the hundreds of parsec scales probed by these observations, or are present with equal frequency in active and inactive galaxies, then the timescale for nuclear activity must be less than the timescale for matter inflow from these spatial scales. If the dynamical timescale on these spatial scales is an approximate measure of the characteristic inflow time, this suggests an upper limit to the lifetime of an individual episode of nuclear activity is several million years."
    },
    "0212/astro-ph0212358_arXiv.txt": {
        "abstract": "We review the ability of redshift surveys to provide constraints on the Dark Energy content of the Universe.  The matter power spectrum and dynamics at the present epoch are nearly `blind' to Dark Energy, but combined with the CMB they can provide a constraint on the Equation of State parameter $w$.  A representative result from the 2dF galaxy redshift survey combined with the CMB is $w <-0.52$ (95\\% CL; with a prior of $w \\ge -1$), consistent with Einstein's Cosmological Constant model ($w=-1$).  More complicated forms of Quintessence (e.g. epoch-dependent $w$ or $w<-1$) are not yet ruled out.  At higher redshifts, the abundance of galaxies and clusters of galaxies, variants of the Alcock-Paczynski curvature test and cross correlation of the CMB with radio sources look potentially promising, but they suffer from degeneracy with other parameters such as the matter density and galaxy biasing. ",
        "introduction": "Recent cosmological observations suggest the existence of a Dark Energy component in the Universe (e.g. Bahcall et al. 1999; Efstathiou et al. 2002).  However, the data still allow room for a complicated cosmic Equation of State $w = P/\\rho$ which may vary with redshift, as proposed by models of `Quintessence' (e.g. Sahni \\& Starobinsky 2000; Peebles \\& Ratra 2002 for a review).  We shall first clarify notation and convention in this field. The case $w=-1$ corresponds to Einstein's Cosmological Constant. Then the only Dark Energy  parameter is the present-epoch Dark Energy density $\\Omega_{Q0}$, which commonly appears in the literature as $\\lambda \\equiv \\Lambda/(3 H_0^2)$ or $\\Omega_\\Lambda$. If a general Equation of State is allowed then one has to solve for both $w(z)$ (parameterized e.g. as $w(z) = w = const.$ or  $w(z) = w_0 + w_1 z$) as well as for $\\Omega_{Q0}$ (see below). In quoting  results from the literature we shall specify which model of Cosmological Constant or Quintessence is assumed and what is the prior on the allowed range of $w$. For example, $w \\geq -1$  is expected for standard minimally coupled scaler fields (e.g. Sahni \\& Starobinsky 2000), although $w < -1$ is allowed in some scenarios (e.g. Melchiorri et al. 2002).   Here we discuss what can (or cannot) be learned from redshift surveys alone and with input from other cosmic probes such as the Cosmic Microwave Background (CMB).  The methods we consider include galaxy clustering and dynamics, the Alcock-Paczynski curvature test, abundance of clusters and galaxies with redshift, and cross correlation of the galaxy surveys with the CMB to detect the integrated Sachs-Wolfe effect. ",
        "conclusions": "The main conclusions of this review are:  \\begin{itemize}  \\item  The present epoch matter power spectrum and dynamics are almost `blind' to a Dark Energy component. This is actually useful as these measurements can be used to deduce the matter density parameter $\\Omega_m$ independent of any assumption  on the nature of the Dark Energy, and then combined with other probes such as the CMB and Supernovae Ia.  \\item Geometrical tests and counts at high redshift are potentially useful, but they suffer from degeneracy with other parameters such as $\\Omega_m$ and biasing.  \\item The current data are consistent with $w=-1$ (i.e. a Cosmological Constant), but other forms of Quintessence $w(z)$ are still possible.  \\end{itemize}"
    },
    "0212/gr-qc0212129_arXiv.txt": {
        "abstract": "We investigate the matching of continuous gravitational wave (CGW) signals in an all sky search with reference to Earth based laser interferometric detectors. We consider the source location as the parameters of the signal manifold and templates corresponding to different source locations. It has been found that the matching of signals from locations in the sky that differ in their co-latitude and longitude by $\\pi$ radians decreases with source frequency. We have also made an analysis with the other parameters affecting the symmetries. We observe that it may not be relevant to take care of the symmetries in the sky locations for the search of CGW from the output of LIGO-I, GEO600 and TAMA detectors.  \\noindent {\\bf Keywords:} gravitational wave -- methods: data analysis -- pulsars: general. ",
        "introduction": "\\par The first generation of kilometer-scale gravitational wave (GW) laser interferometric detectors with sensitivity in the frequency band 10 Hz to few kHz and ultra cryogenic bar detectors sensitive at frequencies around 1 kHz will start collecting data soon. The TAMA 300 (Tsubona 1995) has already done the first large scale data acquisition (Tagoshi et al. 2001), while LIGO (Abramovici  et al. 1992) and GEO600 (Danzmann 1995) have recently carried out their first science observations. VIRGO (Bradaschia et al. 1991) may become operational in couple of years. Also, an eighty meter research interferometer ACIGA (McCleland et al. 2000) near Perth, Australia is under construction, hoping that it may be possible to extend it to multi-kilometer scale in the future.  \\par  At present, majority of searches are focussed in the detection of {\\it chirp\\/} and burst signals. However interest for the search of continuous gravitational wave (CGW) signals from the output of detectors is growing (Jaranwoski, et. al 1998, Jaranowski \\& Kr\\'{o}lak 1999, 2000; Astone et al. 2002; Brady et. al. 1998, Brady \\& Creighton 2000) due to the possibility of enhancing the signal-to-noise ratio (S/N) by square root of observation time $\\sqrt{T_{obs}}$. An optimistic estimate suggests that the Earth based laser interferometric detectors may detect such signals with an observation time of 1-yr.  \\par The strength of the CGW largely depend on the degree of long-lived asymmetry in the source. There are several mechanism for producing such an asymmetry (Pandharipande et al. 1976;  Bonazzola \\& Gourgoulhon 1996; Zimmermann \\& Szedenits 1979; Zimmermann  1980). The estimates of the asymmetry in neutron stars shows that the amplitude of CGW may be $\\le 10^{-25}$. Hence, long integration will be required to get the signature of signal. But this in turn induces several other problems viz. Doppler modulation and non-stationarity of the detector noise. Consequently data analysis becomes more harder. However, Doppler modulation will provide the information of position of the source in the sky.  \\par It has been realized that with present computing power, coherent all-sky full frequency search in the bandwidth (BW) of the detector of a month long data is computationally prohibitive. Some computationally efficient alternative approaches have been suggested viz. {\\it tracking\\/} and {\\it stacking\\/} (Brady and Creighton 2000; Schutz 1998). {\\it Tracking \\/} involves the tracking of lines in the time-frequency plane built from the Fourier transform (FT) of one day long stretches of data while {\\it stacking\\/} involves dividing the data into day long stretches, searching each stretch for signals, and enhancing the detectability by incoherently summing the FT of data stretches. Also, accurate modeling of GW form, optimal data processing and efficient programming are an integral part of all sky-search.  \\par The basic method to analyze the detector output to get the signature of GW signals rely on how efficiently one can Fourier analyze the data. Fourier analysis of the data has the advantage of incorporating the interferometers noise spectral density. The problem for the search of CGW largely depend on how accurately one can take into account the translatory motion of the detector acquired from the motions of the Earth in Solar system barycentre (SSB) frame. It has been shown (Srivastava and Sahay 2002a,b) that amplitude modulation will only redistribute the power of the frequency modulated (FM) signal in five frequency bands $f \\pm 2f_{rot}$, $f$, $f \\pm f_{rot}$, where $f$ and $f_{rot}$ are the frequencies of the FM signal and the rotational frequency of the Earth respectively. Hence it is sufficient to consider only FM signal for the analysis of the matching of signals from different locations in the sky .  \\par The study of the matching of signals from locations in the sky that differ in their co-latitude and longitude by $\\pi$ radians has been made for 1-d data set (Srivastava and Sahay 2002c). They observed symmetries in the sky locations. In this paper we report that the observed symmetries are not generic. We study this problem for longer observation time and the other parameters which affect the symmetries. In the next Section we briefly review the FT of FM signal. In Section 3, using the concept of fitting factor (FF), we investigate the matching of signals in an all sky-search with reference to the Earth based laser interferometric detector by considering the source location as the parameter of signal manifold and templates corresponding to different source locations for longer observation time. We present our conclusions in Section 4. ",
        "conclusions": "In view of blind all sky search for CGW, we have studied the matching of signals from different source locations assuming the noise to be stationary and Gaussian. For fixed $f_o$, we observed that the matching of signals from locations in the sky that differ in their colatitude by $\\pi$ radians is independent of $T_{obs}$, and $\\phi$. However, it falls with $f_o$. But in longitude the matching of signals falls with $T_{obs}$, $\\theta$, $\\phi$, $f_o$. We believe that the matching of signals will increase for the real data. This is due to the resolution provided by the Fast Fourier transform (FFT). However, it may not be relevant to account this symmetries for the search of CGW from the output of LIGO-I, GEO600 and TAMA detectors. \\par This analysis will be more relevant if one performs hierarchical search (Mohanty \\& Dhurandhar 1996; Mohanty 1998). This search is basically a two step search, in first step the detection threshold is kept low and in the second step a higher threshold is used. The higher threshold is used for those templates which exceed the first step threshold. The study of the matching of signals in different source locations will be also relevant for Laser interferometer space antenna (LISA) (Hough J., 1995). The work has been initiated and may be useful for the data analysis of CGW."
    },
    "0212/astro-ph0212502_arXiv.txt": {
        "abstract": "s{ In the last decade our knowledge on galactic black hole systems and in particular on their high energy behavior has considerably improved. I will briefly review here the main results obtained by the high-energy missions SIGMA/GRANAT, Compton-GRO, Beppo-SAX and Rossi-XTE, on these objects and, in particular, I will discuss the spectral shapes observed at energies higher than 30~keV and the detections of high energy features at $>$ 300~keV. Galactic black holes are indeed main targets for the ESA gamma-ray mission, INTEGRAL, to be launched in October 2002, and for the next gamma-ray missions SWIFT, AGILE and GLAST. We expect that a large amount of data will be collected in the next years and the perspectives for the high energy astrophysics of galactic black hole systems look very promising. } ",
        "introduction": " ",
        "conclusions": ""
    },
    "0212/astro-ph0212028_arXiv.txt": {
        "abstract": "{ We examined the maximum bolometric peak luminosities during type~I X-ray bursts from the persistent or transient luminous X-ray sources in globular clusters. We show that for about two thirds of the sources the maximum peak luminosities during photospheric radius expansion X-ray bursts extend to a critical value of 3.79$\\pm$0.15$\\times$10$^{38}$\\,erg\\,s$^{-1}$, assuming the total X-ray burst emission is entirely due to black-body radiation and the recorded maximum luminosity is the actual peak luminosity. This {\\it empirical} critical luminosity is consistent with the Eddington luminosity limit for hydrogen poor material. Since the critical luminosity is more or less always reached during photospheric radius expansion X-ray bursts (except for one source), such bursts may be regarded as empirical standard candles. However, because significant deviations do occur, our standard candle is only accurate to within 15\\%. We re-evaluated the distances to the twelve globular clusters in which the X-ray bursters reside. ",
        "introduction": "\\subsection {Type~I X-ray bursts}  Type~I X-ray bursts (hereafter X-ray bursts, except when needed) in low-mass X-ray binaries are thermonuclear runaways in freshly accreted material on the surface of a neutron star (for a review see Lewin et al.\\ 1993). The freshly accreted matter is compressed and heated, on a time scale of hours to days, to densities and temperatures adequate for thermonuclear ignition. Due to the strong temperature sensitivity of the nuclear reactions, the ignition nearly always results in a runaway process that leads to a sudden, rapid (between about 1 to 10\\,sec) increase in the X-ray flux. These X-ray bursts last generally for tens to hundreds of seconds, and recur with a frequency (partly) set by the supply rate of fresh fuel.  The spectral hardening during the X-ray burst rise and subsequent softening during the decay reflect the heating and subsequent cooling of the neutron star surface. The X-ray spectra during the X-ray bursts are consistent with black-body emission from a compact object with a radius of approximately 10\\,km and temperature between 1 and 2.5\\,keV.  Spikes with similar duration in X-ray light curves have been seen in MXB\\,1730$-$335, also known as `The Rapid Burster' (e.g.\\ Lewin et al.\\ 1993, and references therein; this source also shows type~I X-ray bursts), and GRO\\,J1744$-$28, or `The Bursting Pulsar' (Lewin et al.\\ 1996; Kommers et al.\\ 1997), but these do not show the characteristics of type~I X-ray bursts (in particular, the cooling during the decay). They are thought to be due to sudden accretion events and are referred to as type~II X-ray bursts (Hoffman et al.\\ 1978).  During some type~I 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 atmosphere expands due to radiation pressure. During expansion and subsequent contraction the luminosity is expected to remain almost constant near the Eddington limit. Since the luminosity scales as $L_{\\rm b}\\propto R^2T^4$ for pure black-body radiation, the effective temperature, $T$, drops when the radius of the photosphere, $R$, expands. X-ray bursts with a photospheric radius expansion/contraction phase are therefore 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. The point at which the neutron star atmosphere reaches the surface again (i.e.\\ at the highest effective temperatures) is called `touch-down'. When the expansion is large the temperature may become so low that the peak of the radiation shifts to the UV wavelengths, and no or little X-rays are emitted. Such photospheric radius expansion events (hereafter radius expansion bursts, except when needed) are recognizable by so-called `precursors' in the X-ray light curves followed by a `main' burst (Tawara et al.\\ 1984; Lewin et al.\\ 1984).  We note that, although the X-ray burst spectra seem to be adequately described by black-body emission, the `true' emission from the neutron star and its close environment is expected to be more complex. Electron scattering dominates the opacity in the neutron star atmosphere, leading to deviations from the original Planckian spectrum (van Paradijs 1982; London et al.\\ 1984, 1986; see also Titarchuk 1994; Madej 1997, and references therein). Similarly, electron scattering in a disk corona (e.g.\\ Melia 1987) and/or winds or outflows from the neutron star surface interacting with the accretion disk (e.g.\\ Melia \\&\\ Joss 1985; Stollman \\&\\ van Paradijs 1985) can affect the original X-ray spectrum.  \\subsection{A standard candle?}  Shortly after the discovery of X-ray bursts, van Paradijs (1978) introduced the idea that the average peak luminosity of type I X-ray bursts is a standard candle, presumably the Eddington luminosity limit for the neutron star photosphere. With that assumption, approximate neutron star radii and distances could be derived for a sample of X-ray bursters. Scaling the derived distances so that they are symmetrically distributed around the Galactic center at $\\simeq$9\\,kpc, van Paradijs (1981) found the average peak (bolometric) luminosity to be 3$\\times$10$^{38}$\\,erg\\,s$^{-1}$. Van Paradijs (1981) realized that X-ray bursters residing in globular clusters are important calibrators, since their distances can be determined independently. Their peak luminosity agreed with the above quoted value of the average peak luminosity of Galactic center X-ray bursters. Verbunt et al.\\ (1984) redid this work, based upon a more extensive sample of X-ray bursters. They found the average peak luminosity of X-ray bursters not located in globular clusters to be 3.5$\\pm$1.0$\\times$10$^{38}$\\,erg\\,s$^{-1}$, when taking for the distance to the Galactic center a value of 8.5\\,kpc. For the X-ray bursters in globular clusters a luminosity value of 4.0$\\pm$0.9$\\times$10$^{38}$\\,erg\\,s$^{-1}$ was derived, i.e.\\ consistent with the above value. A `standard candle' value of 3.7$\\times$10$^{38}$\\,erg\\,s$^{-1}$ was adopted. Cominsky (1981), on the other hand, assumed that the black-body radius is the same for all X-ray bursts and subsequently derived source distances and X-ray burst peak luminosities. Again scaling the average source distance to 9\\,kpc, peak luminosities between $\\simeq$10$^{38}$ and $\\simeq$5$\\times$10$^{38}$\\,erg\\,s$^{-1}$ were found.  Van Paradijs (1978, 1981) cautioned, however, that in a few individual sources the scatter in the X-ray burst peak luminosity was appreciable. It was therefore suggested by Lewin (1982) to use instead only the brightest X-ray bursts. In this way the standard candle became $\\gtrsim$5.5$\\times$10$^{38}$\\,erg\\,s$^{-1}$. Of course, the problem arises that one does not know whether the strongest X-ray bursts observed are also the strongest ones possible. It seemed, however, that those X-ray bursts which showed clear photospheric radius expansion do reach a `true' critical luminosity. But whether this critical luminosity is also a standard candle was still questioned (Basinska et al.\\ 1984). If that would turn out to be the case one can in principle determine the distances to sources which show clear radius expansion bursts; for those sources whose X-ray bursts do not show a radius expansion phase we can still determine an upper limit to the distance using the maximum observed peak flux.  The thirteen persistent or transient luminous ($\\gtrsim$10$^{36}$\\,erg\\,s$^{-1}$) X-ray sources in twelve globular clusters are all low-mass X-ray binaries (e.g.\\ Verbunt et al.\\ 1984; Hut et al.\\ 1992; White \\&\\ Angelini 2001; in 't Zand et al.\\ 2001b). All but one (AC\\,211 in NGC\\,7078) have shown type~I X-ray bursts\\footnote{During a recent outburst GRS\\,1747$-$312 in Terzan 6 showed type~I X-ray bursts (in~'t~Zand et al.\\ 2002, in preparation).}, proving they harbour neutron stars. Several of these X-ray bursts exhibited a photospheric radius expansion phase (see, e.g., Lewin et al.\\ 1993, and references therein). We here show that the peak luminosity of radius expansion bursts from about two third of the X-ray bursters in globular clusters do indeed reach an empirical critical luminosity, and may be used as approximate standard candles for other sources. ",
        "conclusions": "\\label{discussion}  We find that the photospheric radius expansion type~I X-ray bursts from bright X-ray sources in globular clusters reach an empirical bolometric peak luminosity (assuming the total X-ray burst emission is entirely due to black-body radiation and the recorded maximum luminosity is the actual peak luminosity) of roughly 3.8$\\times$10$^{38}$\\,erg\\,s$^{-1}$ (excluding 4U\\,1746$-$37, see also below). Because significant deviations from this value occur (e.g., 4U\\,2129+12) this only approximately confirms an earlier suggestion by Basinska et al.\\ (1984) that radius expansion bursts reach a `true' critical luminosity. Our value for the average bolometric peak luminosity is very close to the `adopted standard candle' value of 3.7$\\times$10$^{38}$\\,erg\\,s$^{-1}$ by Verbunt et al.\\ (1984) and comparable to the mean bolometric peak luminosity of 3.0$\\pm$0.5$\\times$10$^{38}$\\,erg\\,s$^{-1}$ by Lewin et al.\\ (1993) of X-ray bursts with photospheric radius expansion from only four globular cluster bursters (i.e., 4U\\,1722$-$30, 4U\\,1820$-$30, A1850$-$08 and 4U\\,2129+12).  To establish whether radius expansion bursts can be used as standard candles, one first has to verify whether the X-ray burst fluxes during every radius expansion burst in an individual source reach a single maximum value. Observations prior to RXTE and BeppoSAX revealed many (i.e., typically more than ten) radius expansion bursts in one globular cluster source, 4U\\,1820$-$30 (Vacca et al.\\ 1986; Haberl et al.\\ 1987; Damen et al.\\ 1990; see Appendix~\\ref{account}), and two sources outside globular cluster sources, MXB\\,1636$-$53 (Inoue et al.\\ 1984; Damen et al.\\ 1990) and MXB\\,1728$-$34 (Basinska et al.\\ 1984). The bolometric peak fluxes were found to be the same within the errors in all three cases.  The BeppoSAX/WFCs observed many radius expansion bursts from the globular cluster sources 4U\\,1820$-$30 and 4U\\,1722$-$30. We find that the bolometric peak fluxes are also the same within the errors in both cases. A similar conclusion was reached from eight radius expansion bursts seen by the BeppoSAX/WFCs from one source outside a globular cluster, XB\\,1812$-$12 (Cocchi et al.\\ 2000a). The more sensitive RXTE/PCA observations of especially MXB\\,1728$-$34 (e.g.\\ van Straaten et al.\\ 2001; Franco 2001) and MXB\\,1636$-$53 show, however, that the bolometric peak fluxes are not exactly constant, with standard deviations of $\\simeq$9.5\\%\\ and $\\simeq$7\\%, respectively (Galloway et al.\\ 2002a; 2002b). Similar studies in other sources but with less radius expansion bursts (between four to eight per source) indicate standard deviations between 7\\%\\ and 15\\%\\ (Muno et al.\\ 2000; Kuulkers et al.\\ 2002; Galloway et al.\\ 2002a; 2002b). In the case of MXB\\,1728$-$34 it was found that the peak fluxes vary systematically on a monthly timescale; taking this secular variation into account the residual scatter is only $\\simeq$3\\%\\ (Galloway et al.\\ 2002a; 2002b).  For the globular cluster bursters MX\\,0513$-$40, XB\\,1745$-$25 and 4U\\,1746$-$37 only two radius expansion bursts were seen from one instrument. In all three cases similar peak fluxes within the errors were reached for both X-ray bursts (see Appendix~\\ref{account}).  We note that two sources outside globular clusters, MXB\\,1659$-$29 (Wijnands et al.\\ 2002) and EXO\\,0748$-$676 (Gottwald et al.\\ 1986), showed very different peak fluxes during different radius expansion bursts. They both are viewed at a high inclination (e.g.\\ Cominsky \\&\\ Wood 1984; Parmar et al.\\ 1986), so that, although the true critical luminosity probably is reached, the X-ray burst emission may be partly obscured by the accretion disk and/or donor star (i.e.\\ the X-ray burst emission is anisotropic). This may lead to different {\\it observed} maximum peak luminosities at different times. Note that this may also affect the peak luminosities during the radius expansion bursts from the globular cluster bursters 4U\\,1746$-$37 and GRS\\,1747$-$312, since they are also relatively high inclination systems (e.g.\\ Sansom et al.\\ 1993; in~'t~Zand et al.\\ 2000; see also Sztajno et al.\\ 1987, who concluded from other grounds that the X-ray burst emission from 4U\\,1746$-$37 is probably anisotropic).  Excluding the high-inclination systems, we conclude that, since multiple radius expansion bursts in a single source reach the same maximum luminosity to within $\\sim$10\\%\\ and since the maximum luminosity is the same within $\\sim$15\\%\\ for different globular cluster sources, the `true' critical luminosity can only be regarded as an approximate empirical standard candle.  In Fig.~7 we show in grey bands the range in the Eddington luminosity limit for matter with cosmic composition (X=0.73; lower band) and hydrogen-poor matter (X=0; upper band) for a (canonical) 1.4\\,M$_{\\odot}$ neutron star with a radius of 10\\,km. The upper and lower boundaries of each band is determined by the Eddington limit observed by an observer far away from the neutron star with very strong photospheric radius expansion and very weak photospheric radius expansion, respectively (van Paradijs 1979; Goldman 1979; see Lewin et al.\\ 1993). Note that an increase in the neutron star mass shifts the band slightly upwards as well as substantially broadens it. A higher than canonical neutron star mass seems to be required in some models explaining the kHz QPO in neutron star low-mass X-ray binaries (see e.g.\\ Stella et al.\\ 1999; Stella 2001; Lamb \\&\\ Miller 2001), and is derived dynamically for the low-mass X-ray binary and X-ray burster Cyg\\,X-2 (Orosz \\&\\ Kuulkers 1999). Our critical luminosity is close to the Eddington luminosity limit for hydrogen-poor matter. Therefore, our critical luminosity is most likely this Eddington limit (see also e.g.\\ van Paradijs 1978).  Previously, `super-Eddington luminosities' were reported for several sources in our sample (e.g.\\ Grindlay et al.\\ 1980; van Paradijs 1981; Inoue et al.\\ 1981; see also Verbunt et al.\\ 1984) and it was referred to as the `super Eddington limit problem' (e.g.\\ Lewin 1985). Winds or outflows from the neutron-star surface set-up by the high radiation pressure were introduced to resolve (part of) this problem (e.g.\\ Melia \\&\\ Joss 1985; Stollman \\&\\ van Paradijs 1985). Our findings do not support super-Eddington luminosities, except possibly for 4U\\,2129+12 (note that 4U\\,2129+12 {\\it is} consistent with the Eddington limit for hydrogen-poor matter if we assume e.g.\\ a somewhat higher neutron star mass, see above). We attribute this mainly to a better knowledge of the distance and their uncertainties.  Five bursters (including GRS\\,1747$-$312) in our sample show bolometric peak X-ray burst luminosities near 1.5$\\times$10$^{38}$\\,erg\\,s$^{-1}$, which is comparable to the Eddington limit of matter with solar composition. Of these, only the sources 4U\\,1746$-$37 and GRS\\,1747$-$312 showed X-ray bursts with evidence for photospheric radius expansion; the maximum peak X-ray burst luminosity for the other sources could have been higher if radius expansion bursts were seen.  If there is a bimodal distribution of the bolometric X-ray burst peak fluxes it is consistent with the thought that the donors in systems with an ultra-short orbital period ($\\lesssim$1\\,hr: MX\\,0513$-$40, H1825$-$331, 4U\\,1820$-$30, A1850$-$08; see Table~\\ref{GC}) do not provide hydrogen, whereas the donors in systems with a longer orbital period (4U\\,1746$-$37, GRS\\,1747$-$312; see Table~\\ref{GC}\\footnote{The 17.1\\,hr orbital period found previously for 4U\\,2129+12 (Ilovaisky et al.\\ 1993) is {\\it not} attributed to the burster but to AC~211 (White \\&\\ Angelini 2001).}) do provide hydrogen-rich material (e.g.\\ Nelson et al.\\ 1986; Verbunt 1987; see also Juett et al.\\ 2001). The short duration of the X-ray bursts from XB\\,1745$-$25 and 4U\\,1820$-$30 ($\\lesssim$25\\,s) is consistent with the fuel being only helium (Bildsten 1995). However, the X-ray bursts with longer duration ($\\sim$mins) from the other systems can not be accomodated this way (e.g.\\ Fujimoto et al.\\ 1981; Bildsten 1998, and references therein), unless the mass accretion rate onto the neutron star is rather low (see Bildsten 1995). On the other hand, we note that the Eddington limit of hydrogen poor matter may still be reached during long X-ray bursts if a hydrogen rich envelope is ejected during the photospheric radius expansion stage of a helium flash (Sugimoto et al.\\ 1984)."
    },
    "0212/astro-ph0212264_arXiv.txt": {
        "abstract": " ",
        "introduction": "Recent high redshift SN-Ia observations suggest that the Universe at present is undergoing accelerated expansion \\cite{snia}. These findings are consistent with the latest precise observations on the anisotropy of the Cosmic Microwave Background Radiation (CMBR) \\cite {cmb} and also with the observations of the Large Scale Structure (LSS) distribution of galactic clusters and superclusters \\cite{conc}. Consequently, modern cosmology seems to have reached a point of concordance, which may be characterized by the following: We seem to live on a spatially flat, homogeneous and isotropic Universe which, at present, is comprised by about 1/3 of pressureless matter (dark matter mostly) and 2/3 by some other substance, with negative pressure, referred to as dark energy. The nature of this dark energy, however, remains elusive.  The above picture is in excellent agreement with the inflationary paradigm, which was initially introduced to solve the horizon and flatness problems of the Standard Hot Big Bang (SHBB) (and some other problems that were thought to be important at the time, such as the monopole problem) \\cite{inf}\\cite{book}. Inflation, in general, predicts a spatially flat Universe and also provides a superhorizon spectrum of curvature perturbations that result in adiabatic density perturbations which can successfully seed the formation of the observed LSS and the CMBR anisotropy. The spectrum of the curvature perturbations is predicted to be very near scale invariance, which agrees remarkably with the latest data. Hence, the inflationary paradigm is now considered by most cosmologists as the necessary extension of the SHBB, in order to form the Standard Model of Cosmology.  The successes of the inflationary paradigm have motivated many authors to consider a similar type of solution to the dark energy problem at present. Thus, it has been suggested that the current accelerated expansion of the Universe is due to a late-time inflationary period driven by the potential density of a scalar field $Q$, called quintessence (the fifth element, added to cold dark matter, hot dark matter (neutrinos), baryons and photons) \\cite{Q}. The aim for introducing quintessence was to avoid resurrecting the embarrassing issue of the cosmological constant $\\Lambda$, which, if called upon to account for the dark energy, would have to be fine-tuned to the incredible level of \\mbox{$\\Lambda^2\\sim 10^{-120}M_P^4$}, where $M_P$ is the Planck mass, i.e. the natural scale for Einstein's $\\Lambda$.  However, it was soon realized that quintessence suffered from its own fine-tuning problems \\cite{KL}. Indeed, in fairly general grounds it can be shown that at present \\mbox{$Q\\sim M_P$} (if originally at zero) with a mass \\mbox{$m_Q\\sim 10^{-33}$eV}, which is very hard to understand in the context of supergravity theories, where we expect the flatness of the potential to be lifted on internal-space distances of the order of $M_P$. In addition, the introduction of yet again another unobserved scalar field (on top of the inflaton field which drives the early Universe inflationary period) seems unappealing. Finally, a rolling scalar field introduces another tuning problem, namely that of its initial conditions.  A compelling way to overcome the difficulties of the quintessence scenario is to link it with the rather successful inflationary paradigm. This is quite natural since both inflation and quintessence are based on the same idea; that the Universe undergoes accelerated expansion when dominated by the potential density of a scalar field, which rolls down its almost flat potential. This unified approach has been named quintessential inflation \\cite{PV} and is attained by identifying $Q$ with the inflaton field $\\phi$. In quintessential inflation the scalar potential of $\\phi$ is such that it causes two phases of accelerated expansion, one at early and the other at late times.  However, the task of formulating such a potential is not easy and this is why not many successful attempts exist in the literature \\cite{PV}\\cite{ng}\\cite{qmodels}\\cite{jose}\\cite{mine2}\\cite{ipl}\\cite{2exp}. Indeed, successful quintessential inflation has to account not only for the requirements of both inflation and quintessence \\cite{yahiro} but also for a number of additional considerations. In particular, the minimum of the potential (taken to be zero, otherwise there is no advantage over the cosmological constant alternative) must not have been reached yet by the rolling scalar field, in order for the residual potential density not to be zero at present. This requirement is typically satisfied by potentials, which have their minimum displaced at infinity, \\mbox{$V(\\phi\\rightarrow\\infty)\\rightarrow 0$}, a feature referred to as ``quintessential tail''. Thus, quintessential inflation is a non-oscillatory inflationary model \\cite{NO}. Another requirement is that of a ``sterile'' inflaton, whose couplings to the Standard Model (SM) particles are strongly suppressed. This is necessary in order to ensure the survival of the inflaton until today, so that it can become quintessence. Thus, in quintessential inflation the inflaton field does not decay at the end of the inflationary period into a thermal bath of SM particles. Instead, the reheating of the Universe is achieved through gravitational particle production during inflation, a process refereed to as gravitational reheating \\cite{grreh}\\cite{JP}. Because gravitational reheating can be rather inefficient, the Universe remains $\\phi$-dominated after the end of inflation, this time by the kinetic energy density of the scalar field. This period, called kination \\cite{JP} (or deflation \\cite{spok}), soon comes to an end and the Universe enters the radiation dominated period of the SHBB. Note, here that a sterile inflaton avoids the danger of violation of the equivalence principle at present, associated with coupled quintessence \\cite{coupled}, where the ultra-light $Q$ corresponds to a long--range force.  In the models \\cite{PV}\\cite{ng}\\cite{qmodels} the plethora of constraints and requirements which are to be satisfied by quintessential inflation is managed through the introduction of a multi-branch scalar potential, that is a potential that changes form while the field moves from the inflationary to the quintessential part of its evolution. This change is either fixed ``by hand'' (such as in the original model \\cite{PV}) or it is due to a potential with different terms that dominate each at a time \\cite{ng} or it is an outcome of a phase transition, arranged through some interaction of the inflaton with some other scalar fields \\cite{qmodels}. Clearly this requires the introduction of a number of mass scales and couplings, which have to be tuned accordingly to achieve the desired results. Thus, in such models it is difficult to dispense with the fine-tuning problems of quintessence. Attempts to design a single-branch potential in \\cite{jose}, which incorporates natural-sized mass scales and couplings have provided with existence proofs, but the the class of potentials presented are rather complicated. This is due to the so--called $\\eta$-problem of quintessential inflation: Namely the fact that it is almost impossible to formulate a successful quintessential inflationary model with an inflationary scale high enough to satisfy the requirements of Big Bang Nucleosynthesis (BBN) but which neither results in strong deviations from scale invariance in the curvature perturbations spectrum, nor does it need to go over to super-Planckian inflationary scale to solve the horizon problem. The $\\eta$-problem is due to the fact that between the inflationary plateau and the quintessential tail there is a difference of over a hundred orders of magnitude. To prepare for such an abysmal ``dive'' the scalar potential cannot help being strongly curved near the end of inflation, which destroys the scale invariance of the curvature perturbations.  It has been thought that this problem is alleviated when considering inflation in the context of brane-cosmology. Indeed, brane-cosmology allows for overdamped steep inflation \\cite{steep}, which dispenses with the need for an inflationary plateau and, therefore, a curved potential seems no longer necessary. However, attempts to use this idea have still encountered difficulties (see for example \\cite{mine2}\\cite{ipl}) and the most promising results were achieved again with a multi-branch potential (a sum of exponential terms) \\cite{2exp}. In this paper we explain why. It seems that, despite the advantages of steep inflation, brane cosmology back-reacts by creating problems in the kination period. Indeed, we will show that the overdamping effect due to the modified dynamics of the Universe, inhibits the efficiency of kination in achieving a small late-time potential density.  Fortunately, there is another solution to the $\\eta$-problem of quintessential inflation. Indeed, we show that the $\\eta$-problem is substantially ameliorated when considering inflation in the context of the curvaton hypothesis \\cite{curv}. As shown recently in \\cite{mine}, the curvaton hypothesis liberates inflationary models from the strains of the so-called {\\sc cobe} constraint, i.e. the requirement that the amplification of the inflaton's quantum fluctuations during inflation should generate a curvature perturbation spectrum with amplitude that matches the observations of the Cosmic Background Explorer ({\\sc cobe}). The curvaton hypothesis attributes the generation of the curvature perturbations to another scalar field, called the curvaton, changing, thus, the {\\sc cobe} constraint into an upper bound. In \\cite{mine} it has been shown that this effect is rather beneficial to many models of inflation well motivated by particle physics. Here, we demonstrate that it may assist also quintessential inflation in overcoming the $\\eta$-problem. This is because, in the context of the curvaton hypothesis, a curved potential does not necessarily destroy the scale-invariance of the curvature perturbation spectrum. Moreover, it may allow for significant reduction of the inflationary scale, which also proves beneficial for quintessential inflation.  The paper is organized as follows. In Section~2 the dynamics of the Universe is briefly layed out both in the case of conventional and also brane cosmology. In Section~3 we look in more detail into the period of kination, which is crucial for quintessential inflation. In Section~4 we discuss the motivation, characteristics and merits of the exponential quintessential tail, which we adopt throughout the paper. In Section~5 we describe the $\\eta$-problem and demonstrate that brane-cosmology cannot overcome it because it inhibits kination. In order to show this we calculate the constraints imposed on quintessential inflation by the BBN and coincidence requirements. We also study the constraints due to the possible overproduction of gravity waves. In Section~6 we present the alternative idea in order to overcome the $\\eta$-problem, namely the curvaton hypothesis. In Section~7 we demonstrate the curvaton liberating effects on a variant of modular inflation in the context of conventional cosmology. We calculate in detail the allowed parameter space and show that all the relevant requirements are met. In Section~8 we investigate the curvaton liberation effects in the case of brane-cosmology, using an exponential potential. We find that successful quintessential inflation is possible in a certain range of values for the brane tension. We carefully calculate the allowed parameter space and show how all the requirements and constraints are satisfied. Finally, in Section~9 we discuss our results and present our conclusions. Throughout the paper we use units such that \\mbox{$c=\\hbar=1$} in which Newton's gravitational constant is \\mbox{$G=M_P^{-2}$}, where \\mbox{$M_P=1.22\\times 10^{19}$GeV}. ",
        "conclusions": "We have investigated the $\\eta$-problem of quintessential inflation model-building. In the context of a potential with an exponential quintessential tail we have shown that brane cosmology inhibits the period of kination due to the extra friction on the roll-down of the scalar field. This counteracts the beneficial effects of steep inflation towards overcoming the $\\eta$-problem. Hence, we pursued a different approach and considered quintessential inflation in the context of the curvaton hypothesis. We showed that the latter substantially ameliorates the $\\eta$-problem in both the cases of conventional and brane cosmology. To demonstrate this we have studied a toy model of what we called modular quintessential inflation in the case of conventional cosmology and the pure exponential potential in the case of brane-cosmology. In both cases we have shown that the available parameter space for the inflationary scale is not large and it is strongly correlated with the reheating efficiency $\\alpha$. Indeed, for a given $V_{\\rm end}$, we have shown that there is only a small window for $\\alpha$, where successful quintessential inflation is possible. This may seem like a fine-tuning problem. However, it simply reflects the necessary tuning for successful coincidence. The required values for $\\alpha$ are not unreasonable and we should point out that there is nothing special about the present time. Any value of $\\alpha$ would cause some brief acceleration period in the late Universe. We just happen to live in this period. These tuning considerations are even more relaxed if one considers the possibility of multiple unfreezings and refreezings of the scalar field, as discussed at the end of Sec.~\\ref{coincbbn}.  In this paper we have considered the intriguing possibility that the scalar field of quintessential inflation (called the `cosmon' by some authors) is a modulus field, possibly associated with the volume of the extra dimensions, such as the geometrical $T$-moduli of weakly coupled heterotic string theory. The modulus is taken to roll down and away from the origin, where it could have been placed by temperature corrections to its potential during a period of thermalization preexisting inflation, if the origin is a point of enhanced symmetry. In this scenario the inflationary expansion begins with a period of thermal inflation followed by fast-roll inflation, as described in \\cite{mine} for modular thermal inflation. In contrast to \\cite{mine} though, we have supposed that the K\\\"{a}hler corrections introduce an exponential slope to the potential over distances comparable to $m_P$ in field space. Thus, the VEV of the modulus is displaced at infinity, while the modulus is stabilized dynamically by being frozen during the later history of the Universe at a non-zero potential density, causing the present accelerated expansion. This way it may be natural to avoid the excessive supergravity corrections that would otherwise increase the present mass of quintessence to unacceptable values. However, it remains to be seen whether this scenario is possible in the context of a realistic string theory.  Turning to brane cosmology we have focused in the much investigated pure exponential potential, which may also be motivated by string theory considerations. In this case there is no preferred starting point for the roll down of the field as long as the inflationary energy scale is kept below the fundamental scale of the theory. We have seen that the parameter space for successful quintessential inflation is somewhat reduced by the negative overdamping effect of brane-cosmology on kination.  Finally, we have studied the effects of gravitational wave generation on quintessential inflationary model-building. We have shown that gravitational waves will not destabilize BBN if the reheating efficiency is \\mbox{$\\alpha>1$}, which may require some tiny, but non-zero coupling of the inflaton with other fields. In the context of the curvaton hypothesis, however, the gravitational wave constraint is ameliorated by the dilution effect of the entropy production due to the curvaton's decay. This may lower the bound on $\\alpha$ below \\mbox{$\\alpha\\sim 0.1$}, which will render gravitational reheating (and a truly sterile inflaton) acceptable. However, a larger $\\alpha$ may be necessary in order to avoid gravitino overproduction: \\mbox{$\\alpha\\gsim 10^2$}. Note, here, that tiny couplings between the inflaton and the SM fields may have beneficiary side effects, such as baryogenesis \\cite{baryon}.  All in all we have shown that the liberating effect of the curvaton hypothesis enables quintessential inflation to overcome its $\\eta$-problem and enlarges the parameter space for successful model-building. Appealing candidates for the quintessential inflaton (or cosmon) may be string--moduli fields.  \\bigskip  \\noindent {\\Large\\bf Acknowledgments}  \\nopagebreak[4]  \\medskip  \\nopagebreak[4]  \\noindent I would like to thank D.H.~Lyth and J.E.~Lidsey for discussions. This work was supported by the E.U.~network program: HPRN-CT00-00152."
    },
    "0212/astro-ph0212578_arXiv.txt": {
        "abstract": "We present an estimation of the point source (PS) catalogue that could be extracted from the forthcoming ESA Planck mission data. We have applied the Spherical Mexican Hat Wavelet (SMHW) to simulated all-sky maps that include CMB, Galactic emission (thermal dust, free-free and synchrotron), thermal Sunyaev-Zel'dovich effect and PS emission, as well as instrumental white noise . This work is an extension of the one presented in Vielva et al. (2001a). We have developed an algorithm focused on a fast local optimal scale determination, that is crucial to achieve a PS catalogue with a large number of detections and a low flux limit. An important effort has been also done to reduce the CPU time processor for spherical harmonic transformation, in order to perform the PS detection in a reasonable time. The presented algorithm is able to provide a PS catalogue above fluxes: 0.48 Jy (857 GHz), 0.49 Jy (545 GHz), 0.18 Jy (353 GHz), 0.12 Jy (217 GHz), 0.13 Jy (143 GHz), 0.16 Jy (100 GHz HFI), 0.19 Jy (100 GHz LFI), 0.24 Jy (70 GHz), 0.25 Jy (44 GHz) and 0.23 Jy (30 GHz). We detect around 27700 PS at the highest frequency Planck channel and 2900 at the 30 GHz one. The completeness level are: 70\\% (857 GHz), 75\\% (545 GHz), 70\\% (353 GHz), 80\\% (217 GHz), 90\\% (143 GHz), 85\\% (100 GHz HFI), 80\\% (100 GHz LFI), 80\\% (70 GHz), 85\\% (44 GHz) and 80\\% (30 GHz). In addition, we can find several PS at different channels, allowing the study of the spectral behaviour and the physical processes acting on them. We also present the basic procedure to apply the method in maps convolved with asymmetric beams. The algorithm takes $\\sim$ 72 hours for the most CPU time demanding channel (857 GHz) in a Compaq HPC320 (Alpha EV68 1 GHz processor) and requires 4 GB of RAM memory; the CPU time  goes as $O(N_{R_o}{N_{\\rm pix}}^{3/2}log(N_{\\rm pix}))$, where $N_{\\rm pix}$ is the number of pixels in the map and $N_{R_o}$ is the number of optimal scales needed. ",
        "introduction": "\\label{intro} \\footnotetext{e-mail: vielva@ifca.unican.es} The study of the Cosmic Microwave Background (CMB) anisotropies is one of the most powerful tools to understand the Universe. The CMB power spectrum analysis provides us information about how the Universe was as early as 300000 years after the Big Bang. Acoustic peaks were predicted to be present in the CMB power spectrum for several cosmological models (see Hu et al. 1997 for a review). Recent experiments like BOOMERanG (Netterfield et al. 2002, Rulh et al. 2002), MAXIMA (Hanany et al. 2000), DASI (Halverson et al. 2002), VSA (Rubi{\\~n}o-Mart{\\'\\i}n et al. 2002), CBI (Mason et al. 2002), ACBAR (Kuo et al. 2002) and Archeops (Benoit et al. 2003) have shown the presence of peaks predicted for those flat models. Moreover, very recently the first detection of the E-mode polarisation in the CMB has been claimed (DASI, Kovac et al. 2002). That detection gives even more support to structure formation models via gravitational instability. However, there are several cosmological parameters that have not been estimated yet, due to the tight experimental requirements needed to put constraints on them. Two ambitious projects have been proposed to measure the CMB anisotropies with enough resolution and sensitivity. One of them is the NASA WMAP mission (Bennet et al. 1996), that was launched in 2001 and which first year data have been released (Bennett et al. 2003 and references therein)    The second one is the most important CMB experiment developed up to date: the ESA Planck mission; it will be launched in 2007 and will provide 10 all-sky maps at 9 different frequencies. Planck mission has two different instruments: the Low Frequency Instrument (LFI, Mandolesi et al. 1998) and the High Frequency Instrument (HFI, Puget et al. 1998). The LFI has a set of radiometers at 30, 44, 70 and 100 GHz; whereas the HFI has bolometers at 100, 143, 217, 353, 545 and 857 GHz. The resolution goes from 5$'$ at high frequencies to 33$'$ at the lowest one. The sensitivity goes from a few $\\mu$K at low frequencies to $\\sim 10$mK at the highest frequency channel. In addition, Planck will generate polarisation maps with good resolution and sensitivity. Thanks to all these properties, Planck data will put constraints for all the cosmological parameters with errors lower than $1\\%$.  Although the Planck frequency range is selected to have a low contribution from additional sources, there are several foregrounds that have an important emission within this frequency range. The cleaning up of the microwave maps is one of the most important challenges in the CMB data analysis. It is necessary --in order to have CMB maps as accurate as possible-- to apply mathematical tools to remove foreground contribution in the microwave sky. Moreover, the better knowledge of these foregrounds is another important goal of the Planck mission. There are several features of the foregrounds (Galactic emissions and extra-galactic sources) that are poorly known (like spatial distribution, spectral indices, new extra-galactic source populations, dust temperature, etc.). With the data provided by Planck, we expect to understand some of them.   During the last years, a large number of mathematical tools have been proposed to perform the \\emph{component separation}: Wiener filter (Tegmark \\& Efstathiou 1996, Bouchet \\& Gispert 1999), Maximum Entropy Methods (MEM, Hobson et al. 1998, Stolyarov et al. 2001), Independent Component Analysis (Baccigalupi et al. 2000, Maino et al. 2002). All these works have achieved good component separations, not only on simulated maps, but also on real data (Barreiro et al. 2003). However, there are some problems with these all-component separation methods. The most general one is related to the philosophy of the methods: the sum of all the recovered components must be equal to the analysed map; in other words, if a component has not been recovered correctly, it is possible that another one has been affected by the unrecovered signal. Fortunately, due to the frequency range that is scanned, the cosmological signal is recovered with good precision, albeit some Galactic emissions are poorly recovered. The second problem we would like to point out is that some components are not well described by the particular assumptions made by the method. All these techniques assume that each component emission can be factorized in both a spatial template and a frequency dependent function. This is only true for the CMB emission and the thermal Sunyaev-Zel'dovich effect (SZ). It is not a bad approximation for the Galactic emission (at least over small and medium size areas). However, it is a really bad approximation for the point source emission, and thus other alternatives must be considered to deal with the point source emission (like modeling it as a noise contribution, Hobson et al. 1999) . In fact, the point source emission is the most problematic one for the all-component separation methods and their contribution to the CMB power spectrum could be important at small scales (Toffolatti et al. 1998).   In order to avoid these problems, other methods have been proposed to remove just one of the microwave sky components. The maximum effort have been done to subtract the emission due to extra-galactic point sources: matched filter (Tegmark \\& Oliveira-Costa 1998, Naselsky et al. 2002), scale-adaptive filter (SAF, Sanz et al. 2001, Herranz et al. 2002a, Chiang et al. 2002) and based on the Mexican Hat Wavelet techniques (MHW, Cay{\\'o}n et al. 2000, Vielva et al. 2001a). Recently, some papers have appeared focused on the SZ detection: SAF (Herranz et al. 2002b, 2002c) and Bayesian approaches (Diego et al. 2002). The results achieved with these methods are promising, since the number of detections and the recovery errors are very good. In addition, the assumptions made in the analysis are very simple. Looking at the previous facts, it is obvious that a combination of both kind of methods (all-component and one-component separation) can achieve better results. For example, the combination of MEM and the MHW (Vielva et al. 2001b) have obtained a very good separation.   \\begin{table} \\begin{center} \\begin{tabular}{|c|c|c|c|c|} \\hline Frequency & FWHM & Pixel size & $N_{\\rm side}$ & $\\sigma_{noise}$ \\\\ (GHz) & (arcmin) & (arcmin) &  & $(10^{-6})$ \\\\ \\hline 857 & 5.0 & 1.72 & 2048 & 19370.15 \\\\ \\hline 545 & 5.0 & 1.72 & 2048 & 426.90 \\\\ \\hline 353 & 5.0 & 1.72 & 2048 & 41.82 \\\\ \\hline 217 & 5.5 & 1.72 & 2048 & 13.76 \\\\ \\hline 143 & 8.0 & 3.44 & 1024 & 4.65 \\\\ \\hline 100 (HFI) & 10.7 & 3.44 & 1024 & 5.29 \\\\ \\hline 100 (LFI) & 10.0 & 3.44 & 1024 & 12.49 \\\\ \\hline 70 & 14.0 & 3.44 & 1024 & 14.66 \\\\ \\hline 44 & 23.0 & 6.87 & 512 & 5.93 \\\\ \\hline 30 & 33.0 & 13.74 & 256 & 3.84 \\\\ \\hline \\end{tabular} \\caption{\\label{planck} Experimental constrains at the 10 Planck channels. The antenna FWHM is given in column 2 for the different frequencies (a Gaussian pattern is assumed). Characteristic pixel sizes are shown in column 3. We show the \\emph{$N_{\\rm side}$} HEALPix parameter in column 4. The fifth column contains information about the instrumental noise level, in $\\Delta T/T$ per pixel.} \\end{center} \\end{table}  In this paper we extend the work in Vielva et al. (2001a). In that paper, the MHW technique was applied to detect point sources in simulated flat sky patches. Now, we apply the spherical generalization of the MHW to detect the point source emission in all-sky maps, following the Planck mission characteristics. We present a Planck point source catalogue that could be achieved with the Spherical Mexican Hat Wavelet (SMHW, Mart{\\'\\i}nez-Gonz{\\'a}lez et al. 2002). The HEALPix scheme (G{\\'o}rski et al. 1999) is used since it is the one expected for the Planck data.  \\begin{table*} \\begin{center} \\begin{tabular}{|c|c|c|c|c|c|c|c|c|} \\hline Frequency & CMB & SZ & TDust     & TDust             & Free-free & Free-free         & Synch.    &  Synch.             \\\\ (GHz)     & ($\\times10^{-5}$) & ($\\times10^{-6}$)   & all-sky   & $|b|>50^{\\circ}$  & all-sky   & $|b|>50^{\\circ}$  & all-sky   & $|b|>50^{\\circ}$  \\\\ \\hline 857       & $4.29$ & $30.10$ & $12.56$             & $0.14$         & $1.75\\times10^{-3}$ & $1.32\\times10^{-5}$ & $3.07\\times10^{-4}$ & $6.79\\times10^{-5}$  \\\\ \\hline 545       & $4.29$ & $15.2$ & $9.40\\times10^{-2}$ & $1.12\\times10^{-3}$ & $4.68\\times10^{-5}$ & $3.75\\times10^{-7}$ & $1.03\\times10^{-5}$ & $2.14\\times10^{-6}$  \\\\ \\hline 353       & $4.29$ & $6.04$ & $4.87\\times10^{-3}$ & $6.06\\times10^{-5}$ & $9.59\\times10^{-6}$ & $8.02\\times10^{-8}$ & $2.60\\times10^{-6}$ & $5.08\\times10^{-7}$  \\\\ \\hline 217       & $4.29$ & $0$                 & $6.12\\times10^{-4}$ & $7.87\\times10^{-6}$ & $6.31\\times10^{-6}$ & $5.46\\times10^{-8}$ & $2.18\\times10^{-6}$ & $3.93\\times10^{-7}$  \\\\ \\hline 143       & $4.26$ & $2.36$ & $1.88\\times10^{-4}$ & $2.46\\times10^{-6}$ & $8.53\\times10^{-6}$ & $7.53\\times10^{-8}$ & $3.16\\times10^{-6}$ & $6.07\\times10^{-7}$ \\\\ \\hline 100       & $4.26$ & $3.41$ & $8.55\\times10^{-5}$ & $1.13\\times10^{-6}$ & $1.43\\times10^{-5}$ & $1.28\\times10^{-7}$ & $7.18\\times10^{-6}$ & $1.14\\times10^{-6}$  \\\\ (HFI, LFI) & & & & & & & & \\\\ \\hline 70        & $4.26$ & $3.96$ & $4.33\\times10^{-5}$ & $5.78\\times10^{-7}$ & $2.72\\times10^{-5}$ & $2.44\\times10^{-7}$ & $1.61\\times10^{-5}$ & $2.42\\times10^{-7}$ \\\\ \\hline 44        & $4.17$ & $3.38$ & $1.92\\times10^{-5}$ & $2.57\\times10^{-7}$ & $6.87\\times10^{-5}$ & $6.21\\times10^{-7}$ & $5.05\\times10^{-5}$ & $7.05\\times10^{-6}$  \\\\ \\hline 30        & $3.99$ & $2.55$ & $1.01\\times10^{-5}$ & $1.36\\times10^{-7}$ & $1.53\\times10^{-4}$ & $1.39\\times10^{-6}$ & $1.35\\times10^{-4}$ & $1.77\\times10^{-5}$   \\\\ \\hline \\end{tabular} \\caption{\\label{rms} We show the RMS values for the map components at the Planck channels. The data correspond to the unconvolved maps are given in $\\Delta T/T$ units. The CMB RMS are presented in the second column; as it is well known, the CMB is frequency independent: the small increase from low to high frequency is due to the different pixelization. In the third column the SZ RMS are shown. The thermal dust, free-free and synchrotron all-sky RMSs values are presented in columns 4, 6 and 8 respectively; whereas in columns 5, 7 and 9 we report the RMS for the same components at Galactic latitudes (absolute value) above $50^{\\circ}$. As the main results of this work are presented for simulations that do not include rotational dust emission, we do not include in the Table the RMS values. We just describe qualitatively its importance: it is the dominant all-sky emission from 30 GHz ($\\approx 10$ times larger) up to 70 GHz ($\\approx 2$ times larger), whereas at Galactic latitudes greater than $50^{\\circ}$ its emission is stronger than the one due to free-free but lower than the synchrotron one.} \\end{center} \\end{table*}  The paper is organized as follows. In Section \\ref{simul} we explain how the simulations have been done and which are the Planck instrumental features. The SMHW is present in Section \\ref{method} as well as the algorithm and the HEALPix implications to the method. The results are given in Section \\ref{results}. In Section \\ref{beams} we introduce the basic procedure to apply the method in maps convolved with asymmetric beams. Finally, the conclusions of this work are presented in Section \\ref{fin}. ",
        "conclusions": "\\label{fin} We have extended the work done in Vielva et al. (2001a) focused on the detection of PS in simulated microwave flat patches of the sky, to the whole celestial sphere. We have done a spherical analysis using the Spherical Mexican Hat Wavelet that is obtained as an stereographic projection of the plane Mexican Hat Wavelet. This extension implies several new steps in the MHW detection method. First of all, we need to estimate the optimal scale at each area of the sky, in order to achieve the maximum number of detections. To perform the optimal scale determination of the spherical maps is a huge task that involves several transformations (from harmonic domain to wavelet space and viceversa) requires a long CPU time and/or storage space. Using the fact that the HEALPix package allows to identify quickly pixels in the sky, we have done plain projections of HEALPix pixels (at $N_{\\rm side}$ = 4) to calculate the optimal scale using Fourier Transform. This approach reduces enormously the CPU time to estimate the optimal scale. We have checked that the error introduced in the optimal scale determination due to the projection is lower than $5\\%$. We have tested that the amplification (flux limit) is greater (lower) than $1\\%$ the one obtained using the exact value of the optimal scale. A big effort has been also done to reduce the CPU time needed for the harmonic transforms. Whereas the SMHW convolution is done in the harmonic domain, the PS detection is done in wavelet space. In addition, three adjacent scales are required for each optimal one, in order to perform a multiscale fit of the wavelet coefficients to better estimate the PS amplitudes. We have solved this problem by modifying some of the HEALPix package codes to perform several harmonic transformations at the same time using pre-computed Legendre polynomials.  Applying the presented algorithm with the detection criterion based on the simulations (see Subsection~\\ref{detect}) that is able to produce PS catalogues with a maximum percentage of spurious detections ($5\\%$) --where spurious means an error larger than $50\\%$--, we are able to recover a PS catalogue for each of the 10 Planck channels (Table \\ref{catalogue}). We can also provide information of complete PS catalogues at different levels ($85\\%$, $90\\%$, $95\\%$ and $99\\%$). These catalogues are in Table \\ref{CatalogueCompleteness}. By subtracting the detected PS from the input PS maps, we can construct PS residual maps that show us how the removed sources are in the high flux tails of the PS distributions: we are just able to detect the brightest sources. However, removing these sources is very important in order to apply all-component separation methods, as MEM or FastICA. These methods assume that the different component emissions can be factorized in a spatial template and a frequency dependent pattern. This is clearly false for the PS emission. To deal with that emission, the all-component separation methods needs to assume the PS signal as a \\emph{noisy contribution}, which for simplicity, is assumed to be Gaussian (different distribution functions for the PS require to add more assumptions to the method). Although the presented residual PS maps are not Gaussian, they are closer to Gaussian than the input PS distribution. In addition, they have RMS values significantly lower than the RMS input maps, what makes the PS emission less important (see Vielva et al. 2001b).  The largest number of detections are at high and intermediate frequencies. At high frequencies the dominant PS populations are due to galaxies whose emission is dominated by dust whereas, at low frequencies, almost all the bright detected sources are flat-spectrum AGN and QSOs. The channels in which the detection result easier are the intermediate ones, in which the Galactic emission is at a minimum. In fact, in the maps where the Galactic emission is relatively more important (at high and low frequencies) the number of local optimal scales needed increases (see Table 3).  By following the detected PS emission in several channels, we are able to provide information about the spectral behaviour of the different PS populations. This is an important result, since, at present, there is not much information about the PS emission at the Planck frequencies (specially at intermediate ones). We are able to follow such emission in several channels, and a few of them along all the Planck frequency range. By fitting the estimated PS amplitudes in small frequency ranges, we can determine sources spectral indices. That is important in order to study the physical processes that produce such emission. We can recover the spectral index with good accuracy as shown in Table \\ref{index}.  We would like to remark that the method is robust in the sense that previous knowledge about the underlaying signals present in the map are not needed: all the information to perform the analysis (the optimal scales) come from the data itself. The only relevant assumption made in the method is that the beam has a Gaussian pattern.  An additional support for the robustness of the method is related to the good detection performance for asymmetric beams, like the ones expected from Planck. The influence of these realistic beams has been tested using the antenna of the Planck satellite with the largest asymmetry (the LFI 28 beam at 30 GHz). Even more, the SMHW can be useful to characterise the beam asymmetry, by permorfing a multi-scale analysis. An specific study concerning the influence on asymmetric beams, non-uniform noise distribution and other systematic effects for the detection of point sources, will be presented in a future work.  Finally, we plan to combine this method with the MEM and FastICA spherical algorithms, in order to perform the all-component separation of the microwave sky."
    },
    "0212/astro-ph0212052_arXiv.txt": {
        "abstract": "This Discussion session focused on star formation within galactic scales.  We attempt to identify the dominant physical processes and parameters that characterize star formation, and to identify key questions that illuminate these phenomena.  The Discussion was delineated by the following cycle of three questions:  (A) Is the top of the {\\sc H\\thinspace ii} LF physically distinct?\\ \\  (B)  How does massive star feedback affect the ISM and star formation?\\ \\ (C) How do ISM properties affect the {\\sc H\\thinspace ii} LF?  Finally, is one of these three questions the fundamental one of the cycle?  Corresponding answers emerged from the Discussion:  (A) The {\\sc H\\thinspace ii} LF to date is a continuous power law;  (B) There are both positive and negative feedback effects, which are poorly understood;  (C) The {\\sc H\\thinspace ii} LF appears remarkably independent of ISM properties.  Therefore, we suggest that the resultant fundamental question is:  {\\bf ``Is the {\\sc H\\thinspace ii} LF and parent stellar cluster membership function universal?''}  This is analogous to the related question of a universal stellar IMF. Understanding the relationship, if any, between the IMF, cluster membership function, and ISM properties may finally lead to a quantitative theory of star formation. ",
        "introduction": "As is suggested by the title of this conference, star formation is one of the fundamental drivers of evolution in the Universe.  It can be considered on scales ranging from individual stars, to starbursts, to the cosmic history of star formation itself.  Roberto Terlevich has spent much of his career considering this entire range of star formation events and interpreting the underlying physical processes that drive the evolution of galaxies and the Universe.  Star forming events are telling us many things, some of which we understand, and many of which we still do not.  In this sense, there are some key works that attempt fundamental interpretations of information from star-forming regions.  Roberto has been a leader in this approach, and he, together with Jorge Melnick, became a front runner when proposing the ``Size -- Velocity Dispersion'' diagramme applied to giant {\\sc H\\thinspace ii} regions (G{\\sc Hii}Rs).  The Terlevich \\& Melnick law (Terlevich \\& Melnick 1981) shows the existence of a correlation between radius and velocity dispersion in G{\\sc Hii}Rs, following a $ R \\sim \\sigma^{2}$ relation, the same as that for globular clusters and elliptical galaxies.  Also the luminosity $L$ and $\\sigma$ are related following the relation, $L\\sim \\sigma^{4}$. Despite the use of the diagrammes as distance indicators, as was their initial attempt, their interpretation of the big star-forming regions was important.  From their findings, Terlevich \\& Melnick proposed that G{\\sc Hii}Rs are virialized systems like globular clusters and elliptical galaxies, and the measured diagnostics indicate the mass of the total system.  Since then there has been much debate on the subject, both on the existence of the correlations themselves and on the interpretation provided.  To reproduce the correlation and to confirm its validity have involved many of us for a number of years.  We have been exploring the definition of the parameters, namely the radius, which is crucial to infer the other two, as well as the shape of the emission lines, which are used to get the $\\sigma$ value (see Melnick et al. 1988; Mu\\~noz-Tu\\~n\\'on 1994; Mu\\~noz-Tu\\~n\\'on et al. 1995 and references therein).  Later on, the Terlevich \\& Melnick law was extended to {\\sc H\\thinspace ii} galaxies (Melnick et al. 1988; Telles \\& Terlevich 1993), and it was demonstrated that under some restriction on surface brightness (Fuentes-Masip et al. 2000) and analysis of the line profiles, the correlations still hold (see Telles et al. 2001; and Telles, this volume). This may indicate that the process of massive cluster formation is similar in both {\\sc H\\thinspace ii} galaxies and G{\\sc Hii}Rs.  The interpretation of the supersonic broadening of emission lines (first reported by Smith \\& Weedman 1970) and its correlation with size is still an issue of debate.  Stellar winds have been blamed for the supersonic width (see, e.g., Chu \\& Kennicutt 1994, debated in Tenorio-Tagle et al. 1996).  Supersonic turbulence is also proposed as a responsible mechanism for the correlations and line broadening, and much work has been carried out both observationally (see the pionering work by Casta\\~neda 1988; and more recently by Joncas 1999, Miesch et al. 1999, and references therein) and theoretically (see V\\'azquez-Semadeni 1999) to address this issue, which remains open.  What is interesting although a bit worrisome, is that the Terlevich \\& Melnick laws are used nowdays to infer physical parameters (namely the mass) of objects at high redshift.  Somehow, and despite the controversies mentioned above, they have been successful in breaking out to become one of the important tools for interpreting star-forming galaxies at high $z$. ",
        "conclusions": "In this discussion session, we continue in the spirit of searching for fundamental interpretations of star-forming regions. Here, we consider star formation within galactic scales.  We attempt to make sense of the various phenomena related to star formation, and we try to identify the dominant processes that affect it. In organizing the topics for the discussion, each of us (CMT \\& MSO) independently compiled a list of our favorite relevant questions.  We found that these issues all appear to fall within the confines of the following three major questions: \\pagebreak  {\\bf \\begin{itemize} \\item A.  Is the top of the {\\sc H\\thinspace ii} LF ({\\sc H\\thinspace ii} LF) physically distinct?  What determines spatial concentration of star formation and starbursts? \\item B.  How does massive star feedback affect the ISM and star formation? \\item C.  How do ISM properties such as metallicity, clumpiness, turbulence, and phase balance affect the {\\sc H\\thinspace ii} region luminosity function and star formation? \\end{itemize} }  These three questions form a cycle:  The {\\sc H\\thinspace ii} LF and especially, the most luminous star-forming regions, determine the nature of massive star feedback; feedback drives properties of the interstellar medium (ISM); and ISM properties presumably must affect the {\\sc H\\thinspace ii} LF. Is one of these three questions, A, B, or C, the fundamental determinant of the other two, and thereby of galactic star formation in general?  How did this seemingly chicken-and-egg cycle originate? Are there actually external relevant issues that break the cycle? The discussion session raised a number of observational and theoretical issues that address these questions.  In \\S 6 below, we nominate our choice for The Fundamental Question."
    },
    "0212/astro-ph0212322_arXiv.txt": {
        "abstract": "We present the results obtained with {\\it Beppo}SAX observations of the three Composite Seyfert/star-forming galaxies: IRAS 20051-1117, IRAS 04392-0123 and IRAS 01072+4954. These sources belong to an enigmatic class of X-ray sources detected in the ROSAT All-Sky Survey (Moran et al. 1996) which is composed of 6 low redshift galaxies. Their optical spectra are dominated by the features of H~II galaxies while their X-ray luminosities ($\\geq$ 10$^{42}$ ergs/s) are typical of Seyfert galaxies. IRAS 20051-1117 shows a 2-10 keV spectrum well described by a power-law with $\\Gamma$ = 1.9 and low intrinsic absorption. This result, the ratio Flux(2-10 keV)/Flux([O~III]$\\lambda$5007) and the significant X-ray variability detected, clearly rule out a Compton thick nature of this source. IRAS 04392-0123 and IRAS 01072+4954, instead, have a {\\it Beppo}SAX flux a factor of $\\sim$ 50-80 smaller than in the previous ROSAT observations, resulting in poor statistics, that prevents detailed modeling. ",
        "introduction": "A large spectroscopic optical survey of bright IRAS and X-ray selected sources from the ROSAT All Sky Survey revealed an enigmatic class of 6\\footnote{The original Moran et al.'s list was composed of 7 objects, but IRAS 10113+1736 is no longer a valid candidate, since infrared and X-ray emission originate from different sources (Condon et al. 1998).} low redshift galaxies with optical spectra dominated by the features of H~II galaxies but X-ray luminosities typical of AGNs, ranging from 1.5 $\\times$ 10$^{42}$ erg/s to 5 $\\times$ 10$^{43}$ erg/s in the ROSAT band (Moran et al. 1996; see Table~\\ref{sample}). These galaxies were named ``Composite''. The diagnostic emission line ratio diagrams (Veilleux \\& Osterbrock 1987) classify these objects as star-forming galaxies. Yet, some of them present [O~III]$\\lambda$5007 lines significantly broader than all other narrow lines in the spectrum and weak and elusive broad H$\\alpha$ wings. Both evidences suggest the presence of a weak or ``obscured'' AGN. An X-ray spectrum is available only for IRAS 00317-2142 observed by ASCA (Georgantopoulos 2000). Its X-ray spectrum can be well reproduced by a single power-law with $\\Gamma$ = 1.7, with low absorption and no detected iron line at 6.4 keV.  Other similar galaxies (i.e. with bright X-ray emission but weak or absent AGN features in the optical band) have been found also in deep ROSAT fields (Boyle et al. 1995, Griffiths et al. 1996) and in the {\\it Chandra} and XMM-{\\it Newton} deep-fields (Rosati et al. 2001, Fiore et al. 2000, Severgnini et al. 2003).  Although these sources usually lack evidence for strong starburst emission, the main problem is the same as for the Moran et al. (1996) sources, i.e., to be able to explain the optical weakness/disappearance of the AGN in these X-ray luminous sources.  It is not likely that the starburst could overpower a Seyfert optical nuclear spectrum, since the starburst component in these objects is not particularly strong (Moran et al. 1996). To explain the absence of optical Seyfert lines we consider three different scenarios: I) the nucleus is heavily obscured and also the NLR is obscured, II) the absorber surrounds almost completely the AGN preventing the UV flux from ionizing the matter otherwise responsible for optical lines and diminish the strength of the Seyfert emission lines, III) The BLRs and the NLRs are absent (Panessa \\& Bassani 2002). Furthermore it is worth note that HST observations of a galaxy previously classified as starburst/H~II on the basis of ground-based observations, have recently revealed a low luminosity Seyfert 2 nucleus hidden by a strong nebular emission from H~II regions near the nucleus (Gezari et al. 2003).  \\begin{table} \\caption{Moran et al. (1996) Composite sample} \\begin{tabular}{lcccc} \\hline \\hline {Name}  & {$N_{HGal}$} & {$F_{0.1-2.4 keV}$} & {$F_{[O~III]}$} & {$F_{IR}$} \\\\ & {$10^{20} cm^{-2}$} & {$10^{-12} (cgs)$} & {$10^{-14} (cgs)$} &  {$10^{-10} (cgs)$} \\\\ \\hline IRAS 00317-2142 & 1.47 & 3.88 & 2.26 & 2.31\\\\ IRAS 00374+5929 & 7.84 & 2.95 & 0.85 & 0.61 \\\\ IRAS 01072+4954 & 13.95 & 1.64 & 6.39 & 0.50\\\\ IRAS 01319-1604 & 1.39 & 1.00 & 1.75 & 1.13\\\\ IRAS 04392-0123 & 5.62 & 2.13 & 1.83 & 0.49 \\\\ IRAS 20051-1117 & 6.53 & 1.30 & 1.27 & 1.20 \\\\ \\hline \\hline \\end{tabular} \\label{sample} \\end{table} ",
        "conclusions": ""
    },
    "0212/astro-ph0212114_arXiv.txt": {
        "abstract": "Despite the interest in dark matter and dark energy, it has never been shown that they are in fact two separate substances. We provide the first strong evidence that they are separate by ruling out a broad class of so-called unified dark matter models that have attracted much recent interest. We find that they produce oscillations or exponential blowup of the matter power spectrum inconsistent with observation. For the particular case of generalized Chaplygin gas models, 99.999\\% of the previously allowed parameter space is excluded, leaving essentially only the standard $\\Lambda$CDM limit allowed. ",
        "introduction": "Despite the broad interest in dark matter and dark energy, their physical properties are still poorly understood. Indeed, it has never even been shown that the two are in fact two separate substances. The goal of this paper is to show that they are.  There is strong evidence from a multitude of observations that there is about six times more cold dark matter (CDM) than baryons in the cosmic matter budget, making up of order 30\\% of critical density \\cite{consistent,efstathiou}. In addition to this clustering dark component, observations of supernovae, the cosmic microwave background fluctuations and galaxy clustering provide mounting evidence of a uniformly distributed dark energy with negative pressure which has come to dominate the universe recently (at redshifts $z\\lesssim 1$) and caused its expansion to accelerate. It currently constitutes about two thirds of the critical density \\cite{consistent,efstathiou,RatraPeebles}.  Although the dark energy can be explained by introducing the cosmological constant ($\\Lambda$) into general relativity (lending the standard model the name $\\Lambda$CDM), this ``solution'' has two severe problems, frequently triggering anthropic explanations and general unhappiness. The first problem is explaining its magnitude, since theoretical predictions for $\\Lambda$ lie many orders of magnitude above the observed value. The second problem is the so-called cosmic coincidence problem: explaining why the three components of the universe (matter, radiation and $\\Lambda$) presently are of similar magnitudes although they all scale differently with the Universe's expansion.  As a response to these problems, much interest has been devoted to models with dynamical vacuum energy, so-called quintessence\\cite{quintessence}. These models typically involve scalar fields with a particular class of potentials, allowing the vacuum energy to become dominant only recently. Although quintessence is the most studied candidate for the dark energy, it generally does not avoid fine tuning in explaining the cosmic coincidence problem. Recently several alternative models have been proposed, such as the condensate models of\\cite{condensates}.   An alternative to quintessence which has attracted great interest lately is the so-called generalized Chaplygin gas (hereafter GCG) \\cite{kamenshchik,bento,fabris,bilic,gorini,dev} (see also the related earlier work of\\cite{barrow}). Rather than fine tuning some potential, the model explains the acceleration of the Universe via an exotic equation of state causing it to act like dark matter at high density and like dark energy at low density. The model is interesting for phenomenological reasons but can be motivated by a brane-world interpretation\\cite{bilic,bento}. An attractive feature of the model is that it can explain both dark energy and dark matter in terms of a single component, and has therefore been referred to as unified dark matter (UDM) or ``quartessence'' \\cite{makler}. (See also \\cite{padmanabhan}.)  This approach has been thoroughly investigated for its impact on the 0th order cosmology, i.e., the cosmic expansion history (quantified by the Hubble parameter $H[z])$ and corresponding spacetime-geometric observables. An interesting range of models was found to be consistent with SN Ia data \\cite{makler} and CMB peak locations \\cite{bento2}.  Some work has also studied constraints from 1st order cosmology (the growth of linear perturbations), finding an interesting range of models to be consistent with cosmic microwave background (CMB) measurements \\cite{carturan}. There is, however, a fatal flaw in UDM models that manifests itself only at recent times and on smaller (Galactic) scales and has therefore not been revealed by these studies. As we will see, this flaw rules out all GCG models except those that are for all practical purposes identical to the usual $\\Lambda$CDM model.  The rest of this {\\it Letter} is organized as follows. In the next section, we review the fundamentals of the GCG model. We then consider in section \\ref{pert} the evolution of density inhomogeneities in the model and use the predicted matter power spectrum to constrain it with observational data. Finally, we describe how the basic flaw that rules out the GCG model is indeed a generic feature of a broad class of unified dark matter models. ",
        "conclusions": "Above we showed that GCG models with $|\\alpha|\\gg 10^{-5}$ are ruled out by observation, since they cause fluctuations or blowup in the matter power spectrum that are not observed. Let us now examine the assumptions that went into this and the broader implications.  First of all, our extremely tight constraints imply that that the narrow range of allowed GCG models will be completely indistinguishable from $\\Lambda$CDM to both to 0th order and at the early times when primary CMB anisotropies are produced. This means that the corresponding standard constraints on cosmological parameters from CMB, SN Ia {\\etc} apply also to the GCG models making the identification $\\Oms=\\Om$, so that there are no interesting degeneracies between $\\alpha$ and other parameters that can significantly widen the allowed $\\alpha$-range. We have therefore used standard constraints $0.2<\\Om<0.4$ for the allowed region in \\fig{exclusion}.  Second, our limiting our constraints to linear scales $k<0.3h/\\Mpc$ was probably overly conservative. As reviewed in \\cite{spacetime,pwindows}, there are quite strong constraints on the linear power spectrum on much smaller scales from weak lensing, from the Lyman $\\alpha$ forest and perhaps even from lensing of halo substructure \\cite{DalalPRL} which if used would tighten our upper limit on $\\alpha$ to $10^{-6}$, $10^{-7}$ and $10^{-10}$, respectively.  Third, we saw that all that really mattered in \\eq{firstorder} as far as the constraints were concerned was the pressure term $(\\lambdac k)^2$. This means that our results apply more generally than merely to the GCG case: {\\it any} unified dark matter model where $p$ is a unique function of $\\rho$ is ruled out if the sound speed is non-negligible, \\ie, if the function $p(\\rho)$ departs substantially from a constant over the range where pressure has an effect --- quantitatively, if $|d\\ln p/d\\ln\\rho|\\simgt 10^{-5}$, again rendering it indistinguishable from standard $\\Lambda$CDM.  In contrast, standard quintessence models have no such problems. Although they typically have high sound speeds causing oscillations as above, this does not prevent the dark matter from clustering since it is a separate dynamic component. Quintessence models would fail as above if there the two components were tightly coupled, and this is effectively what happens with UDM since the two are one and the same substance.  To salvage the UDM idea in some form, its pressure must not be uniquely determined by its density --- not even on subhorizon scales. As worked out in detail by Hu \\cite{HuGDM}, the effective sound speed can under some circumstances differ from the adiabatic sound speed, and it is only the former that must approximately vanish to satisfy our constraints. Although it may be possible to concoct such models, say by introducing another physical field upon which $p$ depends and making it fluctuate in such a way as to cancel the problematic pressure gradients, this would be giving up much of the elegance and simplicity that gave the unified dark matter idea its appeal, essentially substituting one extra field for another.  In conclusion, precision data is gradually allowing us to test rather than assume the physics underlying modern cosmology. We have taken a step in this direction by ruling out a broad class of so-called unified dark matter models. Our results indicate that dark energy is either indistinguishable from a pure cosmological constant or a separate component from the dark matter with a life of its own.   The authors wish to thank Raul Abramo for helpful comments. This work was supported by NSF grants AST-0071213, AST-0134999, AST-0098606 and PHY-0116590, NASA grants NAG5-9194 \\& NAG5-11099, MT and MZ are David and Lucile Packard Foundation fellows and MT is a Cottrell Scholar of Research Corporation."
    },
    "0212/astro-ph0212100_arXiv.txt": {
        "abstract": "We discuss observational constrains coming from supernovae imposed on the behaviour of the Randall-Sundrum models. We test the models using the Perlmutter SN Ia data as well as the new Knop and Tonry/Barris samples. The data indicates that, under the assumption that we admit zero pressure dust matter on the brane, the cosmological constant is still needed to explain current observations. We estimate the  model parameters using the best-fitting procedure and the likelihood method. The observations from supernovae give a large value of the density parameter for brane matter $\\Omega_{\\lambda,0} \\simeq 0.01$ as the best fit. For high redshifts $z > 1.2$, the difference between the brane model and the $\\Lambda$CDM (Perlmutter) model becomes detectable observationally. From the maximum likelihood method we obtained the favored value of $\\Omega_{\\lambda,0} = 0.004 \\pm 0.016$ for $\\Omega_{k,0}=0$ and $\\Omega_{\\mathrm{m},0}=0.3$. This gives the limit $\\Omega_{\\lambda,0} < 0.02$ at $1 \\sigma$ level. While the model with brane effects is preferred by the supernovae type Ia data, the model without brane fluid is still statistically admissible. We also discuss how fit depends on restrictions of the sample, especially with respect to redshift criteria. We also pointed out the property of sensitive dependence of results with respect to choice of ${\\cal M}$ parameter. For comparison the limit on brane effects which comes from CMB anisotropies and BBN is also obtained. The uncertainty in the location of the first peak gives a stronger limit $\\Omega_{\\lambda,0} < 1.0 \\cdot 10^{-12}$, whereas from BBN we obtain that $\\Omega_{\\lambda,0} < 1.0 \\cdot 10^{-27}$. However, both very strict limits are obtained with the assumption that brane effects do not change the physics in the pre-recombination era, while the SN Ia limit is model independent.  We demonstrate that the fit to supernovae data can also be obtained if we admit the phantom matter $p = - (4/3) \\varrho$ on the brane, where this matter mimics the influence of the cosmological constant. We show that phantom matter enlarges the age of the universe on the brane which is demanded in cosmology. Finally, we propose to check for dark radiation and brane tension by the application of the angular diameter of galaxies minimum value test. ",
        "introduction": "The idea of brane universes has originated from Ho\\v{r}ava and Witten \\cite{hw} followed by Randall and Sundrum \\cite{rs1+2}. In the Randall-Sundrum scenarios of the brane-world cosmology the large extra dimensions can solve the mass hierarchy problem of the standard model. In this model the basic cosmological equations are modified by the presence of some extra terms which are derived from the fact that our universe is treated as a three-brane - to which all the gauge interactions are confined - embedded in a five-dimensional, anti-de Sitter space, in (the whole of) which only gravity can propagate. The cosmological evolution of such brane universes has been extensively investigated \\cite{rs1+2,Shiromizu00,Dick01,Dvali00,Deffayet01} (for a review see \\cite{Randall02}). For example, the issues related to the cosmological constant problem \\cite{Tye01}, the cosmological perturbations \\cite{Koyama97,Kodama00}, inflationary solutions \\cite{Dvali01}, homogeneity, and flatness problems \\cite{Brandenberger02,Deffayet02} were discussed.  Brane models admit new parameters which are not present in standard cosmology (brane tension $\\lambda$ and dark radiation ${\\cal U}$). From the astronomical observations of supernovae Ia \\cite{Garnavich98,Perlmutter97,Riess98,Schmidt98,Perlmutter98,Perlmutter99} one knows that the universe is now accelerating and the best-fit model is for the 4-dimensional cosmological constant density parameter $\\Omega_{\\Lambda_4,0}=0.72$ and for the dust density parameter $\\Omega_{m,0}=0.28$ (index \"0\" refers to the present moment of time). In other words, only the exotic (negative pressure) matter in standard cosmology can lead to this global effect. On the other hand, in brane models the $\\varrho^2$ quadratic contribution in the energy density $\\varrho$ even for a small negative pressure, contributes effectively as the positive pressure, and makes brane models less accelerating. In this paper we argue that in order to avoid this problem one requires much stronger negative pressure $p < - \\varrho$ (phantom) matter to be present on the brane (cf. Ref. \\cite{darkenergy,phantcos}).   We concentrate on some observational constraints on the $\\varrho^2$ term, which depending on the type of matter present scales with the cosmic scale factor as $a^{-6\\gamma}$, where $\\gamma$ is the barotropic index in the equation of state $p = (\\gamma -1)\\varrho$ ($p$ - the pressure and $\\varrho$ - the energy density). In fact, the $\\varrho^2$ term scales as $a^{-6}$ (similarly as stiff-fluid in standard cosmology) for dust $\\gamma=1$ matter, as $a^{-8}$ for radiation $\\gamma=4/3$ and as $a^2$ for phantom matter $\\gamma=-1/3$.  In this paper we discuss the observational constraints imposed on the brane-world cosmologies by supernovae type Ia observations. Preliminary results of such analysis were made in \\cite{Godlowski04}. Since we do not know how the exotic physics of brane theory works in the early stages of the Universe (for example,  during structure formation), then all estimations of the brane influence onto the Cosmic Microwave Background (CMB) and Big Bang Nucleosynthesis (BBN) are based on the assumption that this exotic physics does not change physical processes afterwards. The advantage of SNIa data is its independence of the physical processes in the early universe. Therefore, despite the fact that weaker limits from SNIa observations can be obtained, they are in this sense more valuable. We demonstrate that the brane model fits the SN Ia data very well. We also show that these observations indicate a substantial contribution from the brane, although the observational constraints from BBN and CMB restrict this contribution to a much higher degree. However, to derive the limits from BBN and CMB we assume that brane effects did not change physics before the recombination epoch. On the other hand, the less restrictive limit we obtained from SN Ia data is not so strongly dependent on the assumptions of a model.  We argue that the admission of the brane effects for the dust matter on the brane does not fit supernovae data without the presence of the cosmological constant, but we can avoid this problem if we admit phantom matter on the brane. Also, we argue that in the near future, when new supernovae data is available (SNAP3), the hypothesis of important brane effects at present ($\\Omega_{\\lambda,0}>0$) will be tested.  In Section II we present the basic equations of brane cosmology which are studied dynamically in Section III. Sections IV and V are devoted to a redshift-magnitude relation for brane universes and its comparison with supernovae type Ia data. In Section VI we discuss the angular diameter size of a galaxy for brane universes together with the age of the universe problem. In Section VII we obtain the restrictions onto the brane universes from the observations of the Doppler peaks in CMB while in Section VIII we give the restrictions which come from BBN. In Section IX we give our conclusions. ",
        "conclusions": "We have investigated observational constraints for the Randall-Sundrum scenario of brane world cosmology. The motivation to study new observational constraints for some particular scenarios of the brane world cosmology is the demonstration that such models inspired by recent developments in the particle physics can be tested by astronomical observations.  We have shown that this scenario is compatible with the most recent observations of SN Ia. Moreover, we demonstrate that only a new (high $z$) and more precise set of observations can show whether the considered class of models constitute a viable possibility for the description of the present acceleration of the Universe.  The brane model with dust $\\gamma=1$ matter on the brane fits the supernovae Ia data very well (including SN 1997ff at $z \\simeq 1.7$). However, the cosmological constant is still required. We obtain the best-fit non-flat model with $\\Omega_{\\lambda,0} \\simeq 0.01$, $\\Omega_{k,0} \\simeq -0.9$, $\\Omega_{\\mathrm{m},0} = 0.6$, $\\Omega_{\\Lambda,0} =1.3$. For the flat model with $\\Omega_{\\mathrm{m},0}=0.3$ we obtain $\\Omega_{\\lambda_,0}=0.004 \\pm 0.016$ as a best-fit. Whereas the best-fit non-flat model is not realistic (because of large negative curvature), the flat model with the brane for the realistic value of $\\Omega_{\\mathrm{m},0}=0.3$ is in agreement with the SN Ia data.  On the other hand, brane models with phantom $\\gamma = -1/3$ matter on the brane also fit supernovae data and they may {\\it mimic} the contribution from the cosmological constant which is then not required.  We have demonstrated how the values of the estimated parameters depend on the division of the sample on high and low redshifts subsamples. As a results we obtained that although the full data set (Tonry/Barris sample) of SNIa require the cosmological term, the low redshifts $z \\le 0.25$ data set, admits decelerating model without cosmological constant. We have also demonstrated the sensitivity of results with respect to small changes of ${\\cal M}$ parameter.  It is interesting that brane models with dust $\\gamma=1$ matter for high redshifts predict brighter galaxies than the Perlmutter model. Therefore, the difference between the Perlmutter model and the brane model may be detectable for high redshit $z > 1.2$ supernovae.  Let us note that present data suggest that $\\Omega_{\\lambda,0} 1.0 \\simeq 10^{-2}$ but because of the large error of this estimation it is possible that $\\Omega_{\\lambda,0}=0$. However, from future observations of supernovae at high redshift we can expect that errors in estimation of $\\Omega_{\\lambda,0}$ become significantly smaller. At present the brane theory can be neither confirmed nor ruled out. In this way the future SN Ia data should allow to verify the hypothesis that $\\Omega_{\\lambda,0}$ is so large as $\\Omega_{\\lambda,0}\\simeq 0.01$.  We also found the other limits on the value of $\\Omega_{\\lambda,0}$ from the measurements of CMB anisotropies and BBN. We obtain the strongest limits in these cases, namely $\\Omega_{\\lambda,0} < 1.0 \\cdot 10^{-12}$ from CMB and $\\Omega_{\\lambda,0} \\le 1.0 \\cdot 10^{-27}$ from BBN.  However, let us note that because the errors in estimation of $\\Omega_{\\lambda,0}$ from supernovae are so large the results obtained from SN Ia do not contradict those of BBN and CMB.  Of course BBN as well as CMB are very well tested areas of cosmology which do not allow for a significant deviation from the standard model and standard expansion law, except at very early times. Although consistency with BBN and CMB is a crucial issue in brane models, we must remember that in such an approach we assume that brane models does not change the physics in the pre-recombination epochs. On the other hand, the results of analysis SNIa data are independent of the physical processes in the early universe. Therefore weaker limits obtained from SNIa  observations may even be more valuable."
    },
    "0212/astro-ph0212336_arXiv.txt": {
        "abstract": "Over two hundred years ago Capt. James Cook sailed up Whitsunday Passage, just a few miles where we now sit, on a voyage of astronomical observation and discovery that remains an inspiration to us all. Since the prospects of our visiting planets beyond our solar system are slim, we will have to content ourselves with searching for life using remote sensing, not sailing ships. Fortunately, a recently completed NASA study has concluded that a Terrestrial Planet Finder could be launched within a decade to detect terrestrial planets around nearby stars. A visible light coronagraph using an 8-10 m telescope, or an infrared nulling interferometer, operated on either a $\\sim40$ m structure or separated spacecraft, could survey over 150 stars, looking for habitable planets and signs of primitive life. Such a mission, complemented by projects (Kepler and Eddington) that will provide statistical information on the frequency of Earth-sized planets in the habitable zone, will determine key terms in the ``Drake equation'' that describes the number of  intelligent civilizations in the Universe. ",
        "introduction": "This conference takes place in the Whitsunday group of islands discovered on Whitsunday (June 3), 1770, by the English explorer, Capt. James Cook. Surprisingly, the stated goal of Cook's voyage was astronomical in nature. As described in Richard Hough's excellent biography  (1997) of Capt. Cook, the story has echoes in today's searches for other worlds. Johannes Kepler and Edmund Halley (1716) predicted that Venus would traverse across the face of the Sun in 1761 and 1769.\\footnote{Transits occur in pairs separated by roughly a century. The next pair will happen on June 8, 2004, and June 5, 2012. See {\\it The Transit of Venus} (Sellers , 2001).}  These astronomers realized that comparing the timing of the transit from multiple locations would  yield a measurement of cosmological importance, the Earth-Sun distance or Astronomical Unit (AU), then known to no better than a factor of two (Sellers, 2001; and {\\it www.dsellers.demon.co.uk/ venus/ven\\_ch4.htm)}. Just as today's scientists argued that  the Hubble Space Telescope was necessary to determine the Hubble constant, and thus the scale of the Universe, so too did the astronomers of the 18$^{th}$ century argue for an expedition to measure the transit of Venus and thus to determine the scale of the Solar System.  Since the 1761 transit was poorly observed due to bad weather, in 1767 the Royal Society established a committee to ``report on places where it would be advisable to make the observations, the methods to be pursued, and the persons best suited to carry out the work'' of observing the 1769 transit. After a period of deliberation, the Royal Society, acting much like today's National Academy of Science,  forwarded a proposal to King George III calling for an expedition to Tahiti to observe the transit. The proposal offered a  remarkably modern set of arguments (Hough, 1994, p. 44):  \\begin{itemize}  \\item It highlighted the  practical applications of the results, noting that transit measurement would ``contribute greatly to the improvement of astronomy on which Navigation so much depends.''  \\item It noted with concern potential international competition: ``The French, Spaniards, Danes and Swedes are making the proper dispositions for the Observation thereof...The Empress of Russia has given directions for having the same observed...It would cast dishonor [on the British nation] should they neglect to have the correct observations made of this important phenomenon.''  \\item And finally, the proposal included an estimate of the cost of the expedition, exclusive of launch vehicle, that would prove to be an underestimate by a thoroughly modern factor of $\\pi$. ``The expense would amount to {\\it \\char'44}4,000, exclusive of the expense of the ship. The Royal Society is in no condition to defray the expense.''   \\end{itemize}  King George approved the project. The Royal Navy provided the  launch vehicle, the collier {\\it Endeavour}, and Lieutenant James Cook. The Royal Society provided an observer, Charles Green, and the gentleman-naturalist Joseph Banks, who also contributed more than {\\it \\char'44}10,000 of his funds. On August 26, 1768, Cook set off with 94 passengers and crew and supplies for 18 months including 604 gallons of rum and 4 tons of beer. They arrived in Tahiti one year later, six weeks before the transit, and set up an observatory at ``Point Venus.'' The observations were successfully  made on June 3, 1769.  The rest of the voyage was not without event as Cook undertook to fulfill the second (and secret) part of his Admiralty orders, namely to search for the mysterious southern continent. He circumnavigated New Zealand's North and South Islands and explored the East Coast of Australia from Cape Hicks, past Sydney and Botany Bay up to the Great Barrier Reef. On sailing up the Whitsunday Passage, just a few miles from Hamilton Island where we now sit, Cook noted ``This land is diversified by hill and valley, wood and lawn, with a green pleasant experience.'' A few days later, however, he ran aground on the reef, which he termed an ``Insane Labyrinth'', nearly sinking before limping to shore for a month of repairs.\\footnote{The underwater dangers of Australian waters can still catch the unwary mariner, as the near-sinking of Her Majesty's destroyer {\\it HMS Nottingham} after hitting Wolf Rock off the coast of New South Wales during the week of this conference vividly demonstrated.}  Cook returned to England in 1771 with transit data which, when combined with measurements from other sites, led to the determination of the Astronomical Unit to within 3\\% of the modern value, a major advance in observational cosmology (Sellers 2001). ",
        "conclusions": "NASA, together with its potential international partners, has begun to address the challenge of looking for habitable planets and seeking signs of life beyond the Solar System. Captain Cook's voyages remind us that we have asked these questions before and, after much hard work, have been rewarded with new continents to explore.  As Halley said in his paper predicting the 1761 and 1769 transits of Venus (Halley 1716; quoted in Sellers 2001):  \\begin{quote}  {\\it ``We therefore recommend again and again, to the curious investigators of the stars to whom, when our lives are over, these observations are entrusted, that they, mindful of our advice, apply themselves to the undertaking of these observations vigorously. And for them we desire and pray for all good luck, especially that they be not deprived of this coveted spectacle by the unfortunate obscuration of cloudy heavens, and that the immensities of the celestial spheres, compelled to more precise boundaries, may at last yield to their glory and eternal fame.''}  \\end{quote}  Cook's voyage, in the service of science and exploration, resonates with us today as we use transits of extra-solar planets and other techniques to search for new worlds and for life beyond Earth.  That Cook, Banks, and the English society that dispatched them had multiple motivations for  making this voyage --- some noble, some base --- only highlights the similarities with modern exploration where personal ambition, local politics, and practical applications mix with the search for scientific truth. Today's scientists and policy makers would do well to consider the linkages between science and exploration that future voyages of discovery may enable. The technologies needed to look for other habitable worlds are within our grasp. We need only the will (and the funding) to undertake the search."
    },
    "0212/astro-ph0212046_arXiv.txt": {
        "abstract": "Combining X-ray data from the ROSAT PSPC and optical data drawn from the literature, we examine in detail the relationship between the X-ray and optical properties of X-ray bright galaxy groups. We find a relationship between optical luminosity and X-ray temperature consistent with that expected from self-similar scaling of galaxy systems, $L_B \\propto T^{1.6\\pm0.2}$.  The self-similar form and continuity of the $L_B : T$ relation from clusters to groups and the limited scatter seen in this relation, implies that the star formation efficiency is rather similar in all these systems. We find that the bright extended X-ray components associated with many central galaxies in groups appear to be more closely related to the group than the galaxy itself, and we suggest that these are group cooling flows rather than galaxy halos. In addition we find that the optical light in these groups appears to be more centrally concentrated than the light in clusters.  We also use the optical and X-ray data to investigate whether early or late type galaxies are primarily responsible for preheating in groups. Using three different methods, we conclude that spiral galaxies appear to play a comparable role to early types in the preheating of the intragroup medium. This tends to favour models in which the preheating arises primarily from galaxy winds rather than AGN, and implies that spirals have played a significant role in the metal enrichment of the intragroup medium. ",
        "introduction": "\\label{sec:intro}  The environment of a galaxy may have a significant effect on the properties and evolution of the galaxy itself. For example, galaxy clusters are well known to show a relationship between galaxy morphology and density (e.g. \\citealt{dressler80,dressler97}), which suggests some sort of environmental influence on the galaxies. By examining the properties of galaxies as a function of environment it should be possible to gain important insights into the evolution of both galaxies and galaxy systems. Particularly interesting is the group environment, which is typically made up of between three and a few tens of gravitationally bound galaxies. It is important to understand the impact of the group environment on galaxies as the majority of galaxies are found in a group environment \\citep{tully87}, and galaxies in clusters and the clusters themselves will have once been part of groups. In particular, the group environment is likely to have a significant impact on its galaxies as the density and velocities of the member galaxies suggest that mergers and interactions are more common than in clusters or in the field (e.g. \\citealt{mamon92,mamon00}). However, these systems are also small enough that it is possible for the member galaxies to visibly affect the properties of the group as a whole.  X-ray studies of groups of galaxies (e.g. \\citealt{helsdon00,mulchaey98b,mahdavi97,ponman96a,mulchaey96,burns96}) provide information about the group environment and the processes occurring within this environment. Firstly, the presence of diffuse X-ray emission in a group indicates that the group is likely to be a real gravitationally bound object, rather than just a chance superposition of a few galaxies. The X-ray temperature, luminosity and surface brightness profiles provide information about the depth of the potential well and the distribution of mass in these systems. In addition correlations between X-ray parameters can constrain the effects of other non-gravitational processes in these systems. For example the relationship between X-ray luminosity and temperature appears to be steeper in groups than the relation observed in galaxy clusters \\citep{helsdon00,helsdon00b}. This steepening could be the result of the injection of energy into the intragroup medium by either starbursts or AGN, and this effect can be reproduced by theoretical analyses which include the effects of preheating (e.g. \\citealt{balogh99,cavaliere97,cavaliere99}). However, the effects of cooling may also reproduce this steepening if a large fraction of the gas cools out within groups \\citep{bryan00,pearce00,voit01,muanwong01}. Preheating of the gas by the member galaxies is an example of a situation where the galaxies themselves are affecting the evolution and properties of the group. It is also likely that the group environment will play a role in shaping the properties and evolution of the member galaxies.  Little work has been published on the relationship between X-ray properties of groups and optical properties such as total optical light and spiral fraction. Previous work (e.g. \\citealt{mahdavi97,mulchaey96}) has mostly been based on samples which are small or contain optical and X-ray data from a wide variety of sources which may not be directly comparable with one another. In this and a companion paper \\citep{helsdon02b} we aim to look in some detail at the relationship between the X-ray and optical properties of X-ray bright galaxy groups using a consistent approach. In \\citet{helsdon02b} we focus primarily on the morphology-density relation in X-ray bright groups, whilst here we focus more on the overall relationship between X-ray and optical group properties. For the X-ray data we use the 24 galaxy groups analysed in detail by \\citet{helsdon00}. These groups all contain a hot intragroup medium, and thus represent a sample of bound and collapsed groups. We will combine this X-ray data with optical data drawn from the literature, taking care to ensure that the optical data for each group is derived in a way that enables the groups to be directly compared with one another. Examining the optical properties for the groups and their relationship with the X-ray properties will enable us to gain important insights into the evolution and structure of these systems.  This paper is organised as follows. In \\S\\ref{sec:sample} we describe the sample and how we obtain the optical information. We also explain how we attempt to ensure that the optical data for each group is derived in a fair and consistent manner. In \\S\\ref{sec:results} we present the basic analysis of the data, including relations between optical properties, such as total optical luminosity, and X-ray properties such as the group temperature. We also examine the relationship between the central galaxy and the group, and look at the morphological makeup of these systems. In \\S\\ref{sec:dis} we discuss the implications of our results and look in more detail at some aspects such as the origin of preheating in these systems. Our conclusions are summarised in \\S\\ref{sec:conc}. Throughout this paper we use $H_0$=50 km s$^{-1}$ Mpc$^{-1}$. ",
        "conclusions": ""
    },
    "0212/astro-ph0212270_arXiv.txt": {
        "abstract": "The coalescence of massive black hole binaries is one of the main sources of low-frequency gravitational radiation that can be detected by LISA.  When two galaxies containing massive black holes merge, a binary forms at the center of the new galaxy.  We discuss the evolution of the binary after its separation decreases below one parsec.  Whether or not stellar dynamical processes can drive the black holes to coalesce depends on the supply of stars that scatter against the binary.  We discuss various mechanisms by which this supply can be replenished after the loss cone has been depleted. ",
        "introduction": "The prospect that low-frequency gravitational radiation will be detected by LISA has recently energized theoretical inquiries into the formation and the evolution of massive black hole binaries (MBHB).  MBHB are sources of low-frequency gravitational radiation that will provide highest signal/noise for LISA, but the event rate for these sources is much less well known than for the other two principal astrophysical sources, compact binaries and extreme mass ratio inspiral (see contributions by G.~Nelemans and S.~Sigurdsson, this volume).  Astronomical evidence for the existence of black holes with masses $M_{\\rm bh}\\gtrsim10^6 M_\\odot$ in galaxy spheroids with central stellar velocity dispersions $\\sigma\\gtrsim100\\textrm{ km s}^{-1}$ is increasingly compelling; the evidence for black holes with masses $100-10^6 M_\\odot$ in low-dispersion spheroids is still equivocal. When two galaxies merge, a MBHB forms at the center of the new galaxy \\cite{Begelman:80,Roos:81}. There has been considerable interest in determining whether black hole coalescence occurs efficiently following galaxy mergers, since almost all predictions of MBH coalescence event rates equate the galaxy merger rate -- derived from models of structure formation in which galaxies merge hierarchically \\cite{Hae94,MHN01,VHM03,WL03,JB03} -- with the MBH coalescence rate. Detailed estimates of the MBH coalescence rate are fundamentally limited by the resolution of electromagnetic telescopes.  It is therefore expected that MBH coalescences as observed by LISA will facilitate the first definite conclusions regarding galaxy merger rates and MBH demographics at high redshifts \\cite{Hughes:02}.  The mutual gravitational capture of the black holes is facilitated by the dynamical drag imparted to the orbiting MBHs by the overlapping galaxies. The inner parts of elliptical galaxies and many spiral bulges are well described by power-law (``cuspy'') stellar density profiles of the form $\\rho\\sim r^{-\\gamma}$ with $\\gamma\\approx 2$ \\cite{geb96}. High luminosity ellipticals often exhibit a shallowing of the density profile ($\\gamma<2$) near the very centers, where the MBHs are located.  Numerical simulations of galaxy mergers have shown that the black holes remain embedded in their donor cusps throughout the merger \\cite{Milosavljevic:01}. Since the dynamical drag is a function of the combined black hole and stellar cusp mass, which exceeds that of the black hole until the very final stages of the merger, the black hole inspiral and the MBHB formation take place on a dynamical time scale for black holes of comparable mass.  In the unequal mass case, the infall time scale is lengthened proportionally to the black hole mass ratio $M_1/M_2$ ($M_1\\geq M_2$), but for $M_2\\gtrsim 10^6 M_\\odot$ remains short compared to time scale for the subsequent evolution \\cite{Merritt:00}, which is the subject of this contribution. ",
        "conclusions": "We have focused on stellar dynamical mechanisms for extracting energy from a MBHB, but other schemes are possible and even likely. We note a close parallel between the ``final parsec problem'' and the problem of quasar fueling: in both cases, a quantity of mass of order $\\sim 10^8M_\\odot$ must be supplied to the inner parsec in a time much shorter than the age of the universe. Nature clearly accomplishes this in the case of the quasars, probably through gas flows driven by torques from stellar bars \\cite{Shlosman:90}. The same inflow of gas could contribute to the decay of a MBHB, by leading to renewed star formation or rapid accretion of gas \\cite{Armitage:02}.  Similarly, the presence of a third black hole in a galactic nucleus could accelerate the decay by increasing the MBHB's eccentricity through the Kozai mechanism \\cite{Blaes:02}, or by extracting the binary's energy via the triple black hole slingshot interaction \\cite{Valtonen:94}.     \\begin{theacknowledgments} This work was supported by NSF grants AST 00-71099 and by NASA grants NAG5-6037 and NAG5-9046 to DM, and by a Sherman Fairchild Postdoctoral Fellowship to MM. \\end{theacknowledgments}"
    },
    "0212/astro-ph0212450_arXiv.txt": {
        "abstract": "{We carry out a comparison between observations and hydrodynamic simulations of entropy profiles of groups and clusters of galaxies.  We use the Tree+SPH GADGET code to simulate four halos of sizes in the $M_{500}=1.0-16\\times10^{13}h^{-1}M_{\\odot}$\\ range, corresponding to poor groups up to Virgo-like clusters. We concentrate on the effect of introducing radiative cooling, star formation, and a variety of non-gravitational heating schemes on the entropy structure and the stellar fraction. We show that all the simulations result in a correct entropy profile for the Virgo-like cluster. With the heating energy budget of $\\sim 0.7$ keV/particle injected at $z_h=3$, we are also able to reproduce the entropy profiles of groups. We obtain the flat entropy cores as a combined effect of preheating and cooling, while we achieve the high entropy at outskirts by preheating.  The resulting baryon fraction locked into stars is in the 25--30\\% range, compared to 35--40\\% in the case of no preheating. Heating at higher redshift, $z_h=9$, strongly delays the star-formation, but fails to produce a sufficiently high specific entropy. ",
        "introduction": "Clusters of galaxies represent invaluable astrophysical laboratories for the study of the evolution of diffuse cosmic baryons and their interplay with the processes of star formation and galaxy evolution. The observational determination of scaling relations between X-ray properties, such as luminosity $L_X$, gas temperature $T$, and entropy have now demonstrated that this interplay is crucial in establishing the physical properties of the intra-cluster medium (ICM). The slope of the $L_X-T$ relation (e.g.  Markevitch 1998), the amplitude of the mass--temperature relation (e.g. Finoguenov, Reiprich, B\\\"ohringer 2001b) and the gas entropy level (e.g. Ponman, Cannon \\& Navarro 1998; Finoguenov et al. 2002) are all at variance with respect to model predictions based on pure gravitational heating (Kaiser 1986) and call for the need of introducing extra physics to describe the thermodynamics of the ICM (e.g. Evrard \\& Henry 1991). The influence of the energy feedback on the resulting ICM properties caused in turn an interest in using clusters of galaxies as fossil records of the past history of star formation (e.g. Menci \\& Cavaliere 2000; Bower et al. 2001) and the corresponding metal production (e.g. Renzini 1997, Finoguenov, Arnaud \\& David 2001a, Pipino et al. 2002). Following the approach developed in a series of works (e.g. Bower 1997; Ponman, Cannon, Navarro 1999; Tozzi \\& Norman 2001; Borgani et al. 2001), we use the entropy to characterize the properties of the ICM. We adopt the entropy estimator $S={kT / n_e^{2/3}}$ [keV cm$^2$], where $T$ is the gas temperature and $n_e$ is the local electron number density. So far the reproduction of the cluster properties in simulations has generally been restricted to a comparison with general scaling properties of the clusters, such as the $L_X-T$ and the $M-T$ relations (Bialek, Evrard \\& Mohr 2001; Borgani et al. 2002; Muanwong et al. 2002; Dav\\'e, Katz \\& Weinberg 2002; Kay, Thomas \\& Theuns 2002; and references therein).  In this {\\it Letter} we carry out a detailed comparison between the simulations, including the effect of radiative cooling and non-gravitational heating, and X-ray observations of entropy profiles in groups and clusters of galaxies. Our prime goals are to check under which conditions the shape of the observed profiles and their scatter can be reproduced and which is the effect of cooling/heating in determining both the entropy characteristics of the ICM and the amount of gas locked into the cold (stellar) phase. ",
        "conclusions": "Figure \\ref{f:ent} summarizes the comparison between the entropy properties of the ICM for simulated and observed galaxy systems. Panel (a) shows the entropy profiles for observed groups with $M_{500}=(1-3)\\times 10^{13}M_{\\odot}$, compared to the entropy profiles from the simulations of Group-2 and Group-3, panel (b) is for data on groups and poor clusters with $M_{500}=(3-12)\\times 10^{13}M_{\\odot}$ and the simulation of the Group-1, panel (c) is for observed clusters with $M_{500}=(2-10)\\times 10^{14}M_{\\odot}$ and the ``Virgo'' cluster simulation. Panel (d) shows the positions of real and simulated structures on the $(S_{500},M_{500})$ plane (Finoguenov et al. 2002; here $S_{500}$ is the entropy computed at $r_{500}$). Entropy profiles from simulations are plotted down to the radius which contains 100 gas particles. This criterion has been shown by Borgani et al. (2002) to avoid numerical effects and provides results which are stable against numerical resolution.  In panel (a), where we realize the best matching between the masses of simulated and observed structures, the preheating runs with $z_h=3$ reproduce all the entropy characteristics of the observational data.  We note that the S50-3 and K75-3 runs are characterized by comparable amounts of both extra energy and resulting stellar fraction, despite the fact that the two heating recipes spread a different amount of energy to gas at different density.  The differences between the two prescriptions come at radii exceeding $0.5R_{500}$, where the S50-3 approaches the entropy level produced by the gravitational heating case, while in the K75-3 run the entropy keeps increasing parallel to the GH simulations. This behavior could be attributed to the way the gas has been distributed at $z_h=3$, before its accretion onto clusters.  Assigning an equal amount of extra heating to all gas particles is somewhat equivalent to the effect produced by an accretion shock of a fixed strength. Therefore, if self-similarity of the gas holds, the effect of heating will result in a radial behavior of the entropy, which is similar to that of gravitational heating. On the contrary, the S50-3 heating has a marginal effect in increasing the temperature of low-density regions and, therefore, is less efficient in preventing the occurrence of the gravitational shock. As a result, gas undergoes roughly the same amount of gravitational heating as in the self-similar case, whose entropy level is in fact attained in the outer halo regions.  At radii lower than $0.5$--$1\\,R_{500}$ in Fig.1(a-b) the effect of heating is the same in both the K75-3 and the S50-3 runs. According to the energy distribution, one expects the S50-3 to achieve a higher entropy at lower radii. Yet the profiles are nearly the same. This points to an important role of cooling in equating the K75-3 and the S50-3 runs at the center. The difference between the two runs is in the cooling times at high-density regions, which are shorter for K75-3. A similarity in both the entropy profiles and cold fractions indicate that all the high-density gas has been cooled irregardless of the amount of energy put there.  Cooling has been suggested as a mechanism to increase the observed entropy level of the ICM (e.g. Voit et al. 2002, and references therein). Although, by their nature, radiative losses decrease gas entropy, low-entropy particles, which have short cooling time, $t_{\\rm cool}$, are removed from the X-ray emitting phase, thus leaving only high entropy gas, which has longer $t_{\\rm cool}$. However, our results actually show that including star formation and no extra heating produce too steep entropy profiles, an effect which is more evident in the runs of groups. This result suggests that a fair population of low-entropy particles is still left in the diffuse phase during the cooling process. Since cooling of the preheated component is negligible, a combination of preheating and cooling better matches the conditions needed for the model of Voit \\& Bryan (2001) to produce a plateau in the entropy profile: a required separation between the low-entropy gas, which is able to cool and the high entropy gas retained in the hot phase. At outskirts of the groups, however, as shown in Fig.1(d), the observed entropy could only be explained by runs with preheating at $z=3$.  In Fig.\\ref{f:fgas} we outline the difference between simulation runs in accounting at the same time for gas entropy and star fraction.  As expected, the effect of including only gas cooling and star formation is an over-production of the amount of gas locked in the cold phase, while it is only marginal accounting for the observed gas entropy level. The preheating at $z_h=9$ results in a much delayed star-formation, pointing out the need of strongly reducing the energetics of preheating at that epoch. However, even with the amount of energy assigned in our runs, it is not possible to reproduce the high entropy level observed in groups.  \\vspace*{-0.05cm} \\includegraphics[width=8.cm]{plots/ent_fstar_sim.ps} \\figcaption{Simultaneous reproduction of both entropy and amount of star-formation in simulations of clusters and groups. Filled squares are for the runs only including cooling and star formation. Open (filled) circles denote the simulations with preheating at $z_h=9$ ($z_h=3$). The box denotes observational limits discussed in the text. \\label{f:fgas} } \\vspace*{-0.05cm}  Heating at $z_h=3$ performs better in reproducing the observed entropy level, while it still results in some overproduction of the stellar mass fraction. Besides the uncertainties in the observational determination of $f_*$, there is also a number of physical effects, which are not included in our simulations and which should help in further alleviating gas overcooling. For instance, our preheated runs with $z_h=3$ ignore a possible additional feedback from the formation of the first stars (e.g. Madau, Ferrara \\& Rees 2001), which, according to our experience with the $z_h=9$ run, could delay accretion of the gas on the forming halos, both delaying the star-formation (Tornatore et al. 2002) and, in combination with the $z_h=3$ preheating, reducing the fraction of the accreted gas, which could later be available for the star-formation. As a concluding remark, one of the most important results of our analysis is that the entropy profiles in simulations do not depend only on the amount of extra energy, but also on the details of energy ejection, such as the epoch and the spatial distribution. We have demonstrated that the processes of gas accretion, recorded in the gas entropy profiles, also regulate the star-formation.  \\smallskip  {\\it Acknowledgments.} We are very grateful to Volker Springel for providing support in using GADGET, the referee, Mark Voit, for constructive comments, and Trevor Ponman and Paolo Tozzi for useful discussions. The simulations have been run on the IBM-SP3 at the Computing Centers of the Astronomical Observatory in Catania and of the University of Trieste, and on the IBM-SP4 at CINECA in Bologna, under INAF-CINECA grant.  \\vspace*{-0.5cm}"
    },
    "0212/astro-ph0212385_arXiv.txt": {
        "abstract": "We have performed fully 3-D simulations of the collapse of molecular cloud cores which obey the logatropic equation of state.  By following the collapse of these cores from states of near hydrodynamic equilibrium, we are able to produce accretion histories which closely resemble those of observed cores.  The accretion proceeds in four distinct stages: an initial period of very slow accretion; a period of vigorous accretion following the development of a central density singularity, with $\\dot{M} \\propto t^{3}$, as predicted by self-similar models; a period of relatively stable, vigorous accretion;  and finally a gradual decrease in the accretion rate once about 50\\% of the available mass of the molecular cloud core has been accreted.  These results may explain the accretion histories of cores as they pass through the pre-protostellar, Class 0, and Class I stages. ",
        "introduction": "Mounting evidence indicates that massive molecular cloud cores possess properties which are not adequately explained by models employing an isothermal equation of state.  For example, surveys of prestellar cores in a variety of star-forming complexes show that intermediate and high mass cores have nonthermal line widths which are significantly greater than their thermal line widths (\\cite{cm95}).  Models of the collapse of isothermal molecular cloud cores exhibit a number of features which are not borne out by observations, such as a time-invariant rate of accretion onto the central object (\\cite{Shu77}) and highly supersonic infall (\\cite{fc93}).  To address some of the shortcomings of isothermal models, \\cite{mp96} advanced a model for molecular cloud core collapse which accounts for some degree of turbulent support via a phenomenological equation of state.  This logatropic equation of state reads \\\\ $P/P_{c}~=~1~+~A~\\ln(\\rho/\\rho_{c})$ where the subscript ``c'' denotes ``central'' values.  The constant $A$ is an adjustable parameter whose value for a range of molecular cloud core masses was determined empirically by \\cite*{mp96} to be $A = 0.2 \\pm 0.02$.  The logatropic equation of state admits two equilibrium solutions. The features of each solution are developed in \\cite*{mp96}.  The ``singular'' equilibrium solution is analogous to the singular isothermal sphere (SIS) of \\cite*{Shu77}, except that its density profile is $\\rho \\propto r^{-1}$.  The ``nonsingular'' solution is analogous to the Bonnor-Ebert sphere, in that it is a pressure-truncated equilibrium solution with a finite central density and density profile in its outer regions which matches that of the singular solution.  The nonsingular logatrope has the potential to be as useful a model for massive star-forming cores as the Bonnor-Ebert sphere is for low mass star-forming cores (see, for example, \\cite*{alves}, who showed that a Bonnor-Ebert sphere is an excellent fit to the structure of dark cloud B68).  Hence, we undertook a numerical study of the collapse of a nonsingular logatropic sphere from an initially near-equilibrium state until most of its mass had been accreted onto the central condensation. ",
        "conclusions": "The accretion profiles shown in Figure \\ref{f11b} bear several important similarities to what is known observationally about accretion in young stellar objects.  A simple picture of the evolution of an accreting molecular cloud core suggests an accretion history of the following form: accretion proceeds slowly through the pre-protostellar (PPC) phase, increases markedly at the formation of a central source (to power the substantial outflows seen in Class O sources), and finally declines and dissipates as the core passes into the Class I stage.  \\cite*{andre} mark the transition between the Class O and Class I phase at the point where the mass accreted onto the central object is roughly equal to that remaining in the envelope.  Our results show a similar transition at the same point.  As indicated for the 1~M$_{\\odot}$ curve in Figure \\ref{f11b}, there is an initial period of slow accretion, which we liken to the PPC stage.  As soon as the density profile of the core becomes singular ($t=0$), which can be thought of as the moment of protostar formation, the accretion rate increases markedly, launching a phase of vigorous accretion sufficient to power an outflow.  We liken this to the Class 0 phase.  Once half of the available mass has been accreted, the accretion rate declines steadily to near-zero.  We suggest that our simulations of the collapse of logatropic molecular cloud cores bear significant similarities to the observations. They offer an alternative method for modelling the structure of intermediate and higher mass molecular cloud cores.  Cores which exhibit density power-law indices of between 1 and 1.5 (see, for example, \\cite*{vdt}) should be pursued as possible candidate nonsingular logatropes. Our results bolster the argument that models of collapsing protostellar cores \\emph{must} account for some degree of nonthermal support in order to reproduce the observations.   \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{reidma2.eps}} \\caption{Evolution of the radial infall speed, $|v_{r}|$, of a 1~M$_{\\odot}$ nonsingular logatrope preceding the development of the singular density profile.  Adapted from Reid, Pudritz, \\& Wadsley (2002).} \\label{f10} \\end{figure}   \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{reidma3.eps}} \\caption{Time evolution of $M_{acc}(t)$ and $\\dot{M}_{acc}(t)$ for the collapse of nonsingular logatropic spheres with dimensionless radii of $R/r_{o}$ = 1.34 (solid line), 2.21 (dashed line), and 3.26 (dotted line).  The data are scaled such that each core has a central temperature of $T_{c} = 10$~K and surface pressure of $P_{s} = 1.3 \\times 10^{5}$~k$_{B}$, giving them total masses of 1, 2.5, and 5~M$_{\\odot}$.  The vertical dotted lines indicate suggested transitions between the PPC, Class O, and Class I stages for the 1~M$_{\\odot}$ model.  Time $t=0$ marks the moment at which the core's density profile becomes singular.  Adapted from Reid, Pudritz, \\& Wadsley (2002).} \\label{f11b} \\end{figure}"
    },
    "0212/hep-th0212317_arXiv.txt": {
        "abstract": "We study the cosmological evolution of a scalar field that couples to the trace $T=T^{a}_a $ of energy momentum tensor of all the fields (including itself). In the case of a  shallow exponential potential, the presence of  coupling to the trace $T$  in the field equation   makes the energy density of the scalar field decrease faster thereby hastening the commencement of radiation domination. This effect gradually diminishes at later epochs allowing the scalar field to  dominate the energy density again. We interpret this phase as the current epoch of cosmic acceleration with $\\Omega_{\\phi}=0.7$. A variant of this model can lead to accelerated expansion at the present epoch followed by a $a(t)\\propto t^{2/3}$ behaviour as $t\\to \\infty$, making the model free from future event horizon. The main features of the model are  independent of initial conditions. However, fine tuning of parameters is necessary for viable evolution. ",
        "introduction": "The Lagrangian for a system made of the gravitational field, matter sources and the scalar field which couples to the trace of the total energy momentum tensor is derived in \\cite{tpstp88}. (This derivation is briefly summarized in Appendix A for completeness.) The complete action, which takes into account the trace coupling of the scalar field, has the form (see Appendix A): \\begin{eqnarray} S &=& \\nonumber (16 \\pi G)^{-1} \\int{R \\sqrt{-g}} d^4x+{1 \\over 2} \\int { \\beta(\\phi) \\phi^i \\phi_i \\sqrt{-g}}d^4x \\\\ &-&  (8 \\pi G)^{-1} \\int{\\alpha(\\phi) V(\\phi)\\sqrt{-g}}d^4x + S_{\\rm source} \\label{action} \\end{eqnarray} where the coupling to the trace is given by an arbitrary  function $f(\\phi)$ and \\begin{equation} \\alpha(\\phi)={1 \\over {1+4f(\\phi)}}, \\quad \\beta(\\phi)={1 \\over {1+2f(\\phi)}} \\end{equation} The $S_{\\rm source}$ is the remaining source term (radiation, matter ... ), with appropriate coupling to $\\phi$ (see Appendix A). The energy momentum tensor for the field $\\phi$  which arises from the action (\\ref{action}) is given by \\begin{equation} T^{ik}=\\alpha(\\phi)V(\\phi)g^{ik}+\\beta(\\phi) \\left[\\phi^i \\phi^k-{1 \\over 2}g^{ik} \\phi^j \\phi_j \\right] \\label{tikeqn} \\end{equation} We shall assume that scalar field is evolving in an isotropic and homogeneous space time and that $\\phi$ is a function of time alone. The energy density $\\rho_{\\phi}$ and  pressure $p_{\\phi}$ obtained from $T^{ik}$ are \\begin{equation} \\rho_{\\phi}=\\beta(\\phi){ \\dot{\\phi}^2 \\over 2}+\\alpha(\\phi)V(\\phi),~~~~p_{\\phi}=\\beta(\\phi){ \\dot{\\phi}^2 \\over 2}-\\alpha(\\phi)V(\\phi) \\end{equation} The coupling to $T^a_a$  modifies the field energy density $\\rho_{\\phi}$, pressure $p_{\\phi}$ and the equation of state parameter $w_{\\phi}=(p_{\\phi}/\\rho_{\\phi})$.  In a universe containing dust like matter and radiation, the action (\\ref{action}) leads to following  evolution equations for the field: \\begin{eqnarray} \\ddot{\\phi}+{3\\dot{a} \\over a}\\dot{\\phi}&=& \\nonumber{ \\dot{\\phi}}^2 {f_{,\\phi} \\over {1+2f}} +4V(\\phi) {{(1+2f)} \\over {(1+4f)}^2}f_{,\\phi}\\\\ &-&V_{,\\phi}(\\phi) {{(1+2f)} \\over (1+4f)}+{\\rho^i_m \\over a^3}f_{,\\phi}{ (1+2f) \\over {(1+f)^2}} \\label{evoleqn} \\end{eqnarray} (The radiation, being traceless, does not couple to $\\phi$.) The  Friedman equation with the modified energy momentum tensor given by (\\ref{tikeqn}) is \\begin{equation} {\\dot{a}^2(t) \\over a^2(t)}={{8 \\pi G} \\over 3}\\left[{\\dot{\\phi}^2 \\over {2(1+2f)}} +{ V(\\phi) \\over {1+4f}}+{ \\rho_b(a) } \\right] \\label{friedeqn} \\end{equation} where  background energy density due to radiation and matter is given by  \\begin{equation} \\rho_b(a)={ \\rho^i_R \\over a^4}+{\\rho^i_m \\over {a^3(1+f)}} \\end{equation} The simplest form of coupling which we shall adopt is   the linear one \\begin{equation} f(\\phi)=g_c( \\phi/M_p) \\end{equation} where $  M_p =\\sqrt{1 / (8 \\pi G)}$ is the reduced Planck's mass and $g_c$ is a coupling constant. Since  equations (\\ref{evoleqn}) and (\\ref{friedeqn}) reduces to the standard form for  $g_c=0$, this parameter  gives the strength of the forcing term. Equation (\\ref{evoleqn}) describes the evolution of the scalar field in the expanding universe under the influence of an external ``force'', which depends upon the field and its kinetic energy $(1/2)\\dot{\\phi}^2$. The influence of this term on the dynamics of the scalar field is accumulative.   In the absence of matter, the conservation equation formally equivalent to (\\ref{evoleqn}), has the usual form \\begin{equation} \\dot{\\rho_{\\phi}}+3H(\\rho_{\\phi}+p_{\\phi})=0 \\label{conseqn} \\end{equation} and the evolution of energy density is given by \\begin{equation} \\rho_{\\phi}=\\rho_{0\\phi}e^{-\\int{6\\left (1-\\zeta(a)\\right){da \\over a}}} \\end{equation} with $$ \\zeta(a)={1 \\over {(K_e/ P_e) +1}}$$ where the ratio of effective kinetic to potential energy$ (K_e/P_e)$ is given by \\begin{equation} {K_e \\over P_e}={\\beta(\\phi) \\over \\alpha(\\phi)}{\\dot{\\phi}^2 \\over {2 V(\\phi)}} \\end{equation} Interaction modifies the kinetic as well as the potential energy through $ \\alpha(\\phi)$ and $\\beta(\\phi)$, however, $ \\rho_{\\phi}$ bears the same relation to their ratio as in the case of standard scalar field cosmology. Equation  (\\ref{conseqn}), of course,  is not valid in presence of matter; neverthless numerical work shows that the ratio $(K_e / P_e)$ is still  a good parameter which influences the dynamics. The evolution of potential to kinetic energy ratio plays a significant role in  the growth or decay of the energy density $\\rho_\\phi$ at a given epoch and will be crucial in the following discussion.  Let us next address the question of choice of potential of the scalar field which would lead to a viable cosmological model. The obvious restriction on the evolution  is that, starting  from Planck's time,  the scalar field should survive till today (to account for the observed late time accelerated expansion) without interfering with the nucleosynthesis of the standard model.   Exponential potentials  lead to solutions which are the attractors of evolution equations when $g_c=0$ and provide backdrop for the understanding of dynamics in our case. A standard model with \\begin{equation} V(\\phi)=V_0 \\exp\\left[-\\lambda{ \\phi \\over M_p}\\right] \\label{exppot} \\end{equation} (and $g_c=0$) is well studied in literature and can be divided into two categories: (a) The exponential potentials for which $\\rho_{\\phi}$ scales slower than the background density $\\rho_b=(1/a^n)$ which translates to $\\lambda< \\sqrt n$. (b) Potentials for which scalar field energy density scales faster than the $\\rho_b$, i.e. $\\lambda>\\sqrt n$. In the first case if the background energy density was sub-dominant, it would become more sub-dominant and radiation domination will never occur. In the second case there exists a scaling solution which  mimics the background energy density that is dominant, with \\begin{equation} \\Omega_{\\phi}={\\rho_{\\phi} \\over {\\rho_{\\phi}+\\rho_b}}={n \\over \\lambda^2}; \\quad \\rho_{\\phi}\\propto \\frac{1}{a^n} \\end{equation}  The discussion of this model with $g_c\\ne 0$ (so that the coupling to the trace of the stress tensor is switched on) is given in Appendix B. It turns out that this model is not as attractive as another variant of the potential with $V(\\phi)=V_0 e^{\\lambda{ \\phi \\over M_p}}$ which we shall concentrate on next.    \\begin{figure}[p] \\resizebox{3.0in}{!}{\\includegraphics{fig1.eps}} \\caption{ The ratio of  kinetic  to potential energy of scalar field is depicted for the exponential potential (\\ref{pottw+} ) with $\\lambda=5.2$ , $V_0^{1/4}\\simeq 2.8\\times 10^{-30} M_p$ with two different values for coupling (chosen for illustration): $g_c=0.04$ (dashed line) and $ g_c=0.041$(solid line).The kinetic energy ( as compared to the potential energy ) is seen to build up fast to a large value forcing $\\rho_{\\phi}$ into the kinetic regime which otherwise would scale slowly with the scale factor ($\\rho_{\\phi} \\propto a^{-\\lambda^2}$) for the model with the potential in (\\ref{pottw+}). } \\label{figkp+} \\end{figure}    \\begin{figure}[p] \\resizebox{3.0in}{!}{\\includegraphics{fig2.eps}} \\caption{The energy density is plotted against the scale factor:  solid  line corresponds to $\\rho_{\\phi}$ for $g_c=5.9$ in case of the model described by  (\\ref{pottw+}) with $\\lambda =0.52$ and $V_0^{1/4}\\simeq 2.8\\times 10^{-30} M_p$. The dashed and dotted lines corresponds to energy densities of matter and radiation respectively. The  large value of the scalar field kinetic energy (relative to the potential energy ) induced by the forcing term in the evolution equation  ensures that the $\\rho_\\phi$ drops below $\\rho_{b}$, after which  $\\rho_{\\phi}$  remains approximately constant for a period of time. At the end of this phase, the scalar field  tracks the background energy density before becoming dominant and driving the current accelerated expansion of the universe.  } \\label{figden1} \\end{figure}    \\begin{figure}[p] \\resizebox{3.0in}{!}{\\includegraphics{fig3.eps}} \\caption{Evolution of the kinetic to potential energy ratio of scalar field for the model described in figure (\\ref{figden1}) starting from the locking regime onwards. At the end of locking period, the ratio fluctuates about its background value and then nearly tracks it slowly decreasing towards its attractor value. The ratio does not lock itself to its attractor value and approaches zero where it stays asymptotically. The vertical dashed line marks the current epoch.} \\label{figkp1+} \\end{figure}  \\begin{figure}[p] \\resizebox{3.0in}{!}{\\includegraphics{fig4.eps}} \\caption{Evolution of dimensionless density parameter $\\Omega$ is shown as a function of the scale factor  for the model described in figure (\\ref{figden1}) for: (i) the scalar field (solid line), (ii) radiation (dotted line) and matter (dashed line). Late time behavior of  the scalar field leads to  the present day value of $\\Omega_{\\phi}=0.7$ and $\\Omega_{m}=0.3$.} \\label{figomega} \\end{figure}   \\begin{figure}[p] \\resizebox{3.0in}{!}{\\includegraphics{fig5.eps}} \\caption{The  equation of state parameter of scalar field for the exponential potential (\\ref{pottw+}) showing different initial conditions  converge to attractor solution at late times. The horizontal dashed line depicts the $w_b=1/3=w_R$. The solid and dashed lines start with widely different initial conditions but lead to similar behaviour at late times.} \\label{figoustat} \\end{figure}   \\begin{figure}[p] \\resizebox{3.0in}{!}{\\includegraphics{fig6.eps}} \\caption{Evolution of equation of state parameter  of scalar field.  The coupling to the trace $T$  brings scalar field (which was slowly rolling down the shallow potential) into kinetic regime. The scalar field stays in this regime for a long time before $\\rho_\\phi$ grows towards $\\rho_{\\rm b}$. During this phase, $w_\\phi$ gets locked to $w_{\\phi}=-1$ ( locking regime). After the locking period ends, $\\rho_{\\phi}$ tracks the background energy density slowly approaching it from below with equation of state $w_{\\phi} \\simeq 1/3$. At late times $\\rho_{\\phi}$ overtakes the background and becomes dominant to account for the current epoch of cosmic acceleration. During this time $w_{\\phi}$ decreases and heads towards its inflationary attractor value ($ w_{\\phi}=(\\lambda^2-3)/3)$). However, before it could settle there, the slowly moving field $\\phi$ towards the origin gets trapped near the barrier formed there due to coupling to the trace $T$ . The equation of state locks to $-1$  permanently. The vertical dashed line marks the present epoch.} \\label{figeqstat} \\end{figure}   \\begin{figure}[p] \\resizebox{3.0in}{!}{\\includegraphics{fig7.eps}} \\caption{The evolution of energy density in the field is shown for two values of $g_c=0.9$(solid line) and  $g_c=0.07$(dotted line) for the model described (\\ref{pottw+}) with $\\lambda=0.52$. The radiation energy density is depicted by the dashed line. The second dashed line below shows the(non-coupled) inflationary behavior of $\\rho_{\\phi}$($ \\rho_{\\phi} \\propto a^{-\\lambda^2}$). The decay of $\\rho_\\phi$ slows down earlier for smaller value of the coupling constant $g_c$ and approaches the inflationary attractor behavior after overtaking   the background(radiation) energy density. After staying with the attractor value for quite some time, $\\rho_{\\phi}$ becomes constant.} \\label{figdeng1g2} \\end{figure}  \\begin{figure}[p] \\resizebox{3.0in}{!}{\\includegraphics{fig8.eps}} \\caption{The effective potential in equation (\\ref{veffpot}) for $\\lambda =0.52, g_c=5.9$. This  effective potential becomes repulsive near the origin when $(\\phi_{cr}/M_p)=1.88$ followed by an infinite barrier at $(\\phi^b/M_p)=-0.042$ which shields the singularity situated at $(\\phi^s/M_p)=-0.084$.} \\label{figeffpot} \\end{figure} \\begin{figure}[p] \\resizebox{3.0in}{!}{\\includegraphics{fig9.eps}} \\caption{Late time evolution of the scalar field (solid line). After the current epoch, the scalar field passes across the critical point, moves towards the barrier and gets reflected in the effective potential shown in the previous figure. It makes few oscillations and settles at the critical point ($ \\phi_{cr}/M_p) = 1.88$. The dashed line depicts the inflationary attractor behaviour of the scalar field in case of the exponential potential (\\ref{pottw+}) with $\\lambda=0.52$ ($\\phi/M_p=-\\lambda \\ln(a/a_i)+constant$) and the dotted line corresponds to $(\\phi_{cr}/M_p)=1.88$.} \\label{figft} \\end{figure} ",
        "conclusions": "We have examined the dynamics of a scalar field  which couples to the trace of all the field including itself in FRW background. The coupling to the trace $T$  of the field modifies the energy momentum tensor and induces a forcing term in the field evolution equation. The forcing term acts against the Hubble damping changing the evolution  of field energy density significantly. In case of a steep potential, it amounts to increasing  the steepness of the potential. In case of the shallow potentials in which the field rolls slowly and mimics cosmological constant, the coupling to the trace $T$    plays an important role. It forces the scalar field into kinetic regime allowing the  radiation domination to commence. We have discussed a model in which the effect of the back reaction gradually diminishes allowing the scalar field to catch up with the inflationary attractor regime at late times, thereby accounting for the observed cosmic acceleration with $\\Omega_{\\phi}=0.7$. We have also presented a  variant of the model in which the scalar field plays the role of quintessence at the current epoch and mimics  dust like dark matter in future which eliminates the future event horizon. The general features of the model are shown to be independent of the initial conditions. However,  fine tuning of some of the parameters  is necessary to achieve viable evolution.   \\appendix"
    },
    "0212/hep-th0212067_arXiv.txt": {
        "abstract": "We consider a cosmological brane moving in a static five-dimensional bulk spacetime endowed with a scalar field whose potential is exponential. After studying various cosmological behaviours for the homogeneous background, we investigate the fluctuations of the brane that leave spacetime unaffected.  A single mode embodies these fluctuations and obeys a  wave equation which we study for bouncing and ever-expanding branes. ",
        "introduction": "In the last few years, a lot of efforts have been devoted to  the investigation of  the braneworld picture, whereby our accessible universe is a three-dimensional submanifold, or three-brane, embedded  in a higher-dimensional manifold.  The cosmological consequences of this idea are of particular interest since new effects can be  anticipated  in the very early universe where the physical conditions are very different from those of the present universe. Although many scenarios exist in the literature, most  models   of  brane cosmology focus, like  the present work, on a {\\it self-gravitating} brane-universe embedded in a {\\it five-dimensional} bulk spacetime, so that the brane world-sheet is of codimension one and subject to the standard junction conditions for a thin wall in general relativity. As usually assumed, we will take a $Z_2$ symmetric bulk, which means that the two sides of the brane are  mirror-symmetric  with respect to the brane.  The simplest models of  brane cosmology (see \\cite{l02} for a recent review) assume an empty bulk with a cosmological constant. The latter (with a negative sign) is necessary  in order to recover a standard cosmological evolution at late times and in particular to account, via nucleosynthesis, for the abundances of light elements. In the early universe, however, the evolution deviates from standard cosmology.  Although very useful for some specific features of brane cosmology, an  empty bulk might be too naive for a realistic description of the early universe. For example, in the five-dimensional version of M-theory, a scalar field, corresponding to the volume  of the Calabi-Yau compactification manifold, is present in the five-dimensional bulk \\cite{low}. It is thus relevant  to investigate brane cosmology with a bulk scalar field \\cite{cr99}-\\cite{ftw01}, which might be also useful, in the case of two-brane models, to stabilize the radion \\cite {gold}.  In the present work, we consider a five-dimensional model  where the bulk contains   a scalar field with an exponential potential and a three-brane with a cosmological perfect fluid conformally  coupled, via the bulk scalar field, to the induced metric. We restrict our attention to very specific bulk spacetimes with the usual cosmological symmetries, i.e. homogeneity and isotropy along the three ordinary spatial dimensions, that are also {\\it static}   in the two-dimensional subspaces spanned by  time and the extra dimension. We use  {\\it static} here in a generalized sense, where the orbits associated with a Killing vector  can be not only time-like but also space-like.  We first assume that the brane is perfectly homogeneous and study the corresponding background cosmologies  associated with the motion of the brane in such static bulk spacetimes. We thus generalize the results of Chamblin and Reall \\cite{cr99}, restricted to the case of a domain wall, to any equation of state for the brane matter.  We then allow the brane to fluctuate but we impose that these fluctuations are such that {\\it the bulk spacetime is left unperturbed}. In other words, we investigate only the fluctuation mode, which one can call the {\\it intrinsic mode}, that is not coupled to the gravitational radiation, i.e. to the bulk perturbations. We show that this mode obeys  a wave equation, which can be written in a familiar form. We analyse the evolution of the brane fluctuation depending on the various background cosmologies. In some sense, our approach is reminiscent of former studies \\cite{gar,guv} (see also \\cite{ishi} for recent developments) of perturbed test branes where the brane deformation is described by a scalar field obeying a Klein-Gordon equation. In our case the self-gravity is included by adjusting adequately the matter perturbations on the brane. A similar analysis has also been carried out in \\cite{steer} within the context of mirage cosmology where the gravitational back-reaction is neglected.  The plan of our paper is the following. In the next section, we present the framework and consider some background homogeneous solutions. In the third section, we derive the equation of motion for  the brane fluctuations. In section 4, we analyse this wave equation for  the background cosmologies discussed in section 2. Finally, we conclude in the last section. ",
        "conclusions": "We have investigated the fluctuations of a moving brane in a dilatonic background. These fluctuations are represented by a scalar mode on the brane corresponding to ripples along the normal direction to the brane. As the brane fluctuates, it induces metric fluctuations, in particular we have found that the induced metric appears naturally in the longitudinal gauge with two unequal Bardeen potentials $\\Phi$ and $\\Psi$. The fact that these potentials are not equal springs from the presence of anisotropic stress on the brane. For a fixed equation of state for the matter content on the brane, for instance cold dark matter, we find that the two Bardeen potentials are proportional. As such this implies that a single gauge invariant observable $\\Phi$ characterizes the brane fluctuations.  Our approach differs from the projective approach \\cite{mw00,mb00} in as much as we have not considered the perturbed Einstein equations on the brane. This allows us to free ourselves from the thorny problem of the projected Weyl tensor on the brane.  We have focused on the motion and fluctuations of branes in a particular class of dilatonic backgrounds. These backgrounds correspond to an exponential potential and an exponential coupling of the bulk scalar  field to the brane. The motion of the brane is either of the bouncing type  or  the ever-expanding form. In the bouncing case we find that the brane cannot escape towards infinity, it is bound to a singularity which is either null or time-like. In the time-like case, i.e. when it appears at a  finite distance in conformal coordinates, the fluctuations of the brane are unbounded implying that the brane is ripped by the strong gravity around  the singularity. In the null case, i.e. when the singularity is at conformal infinity, the fluctuations oscillate in a bounded manner while converging to the singularity. The bouncing case is equivalent to the behaviour of a brane in a supergravity background. As we only consider intrinsic fluctuations of the brane in an unperturbed bulk, this corresponds to a situation where supersymmetry is preserved by the bulk while broken by the brane motion. Therefore the bouncing brane fluctuations correspond to fluctuations of a non-BPS brane embedded in a supergravity background.  In the ever-expanding scenario, we can distinguish two possibilities. The brane can escape to infinity with a scale factor which is either expanding in a decelerating manner or accelerating, i.e. corresponding to an inflationary era of the power law type. In the decelerating case, the brane eventually oscillates forever. In the inflationary case, the brane is such that any fluctuation of a giving length scales oscillates until it freezes in while passing through the horizon. Of course this scenario is reminiscent of  four-dimensional inflation modelled with a scalar field. Here the features of inflation, i.e. the relationship between the power spectrum and the scale factor, differ from the four dimensional case. This is an interesting observation as it leads to a new twist in the building of inflationary models. One might hope that alternative scenarios to four-dimensional inflation may emerge from five dimensional brane models and their fluctuations."
    },
    "0212/astro-ph0212317_arXiv.txt": {
        "abstract": "We present here the first results of two-dimensional hydrodynamical simulations of the neutrino-heating phase in the collapsed core of a 15$\\,M_{\\odot}$ star, where the neutrino transport is treated with a variable Eddington factor method for solving the Boltzmann transport equation, and the neutrino interactions include nucleon-nucleon bremsstrahlung, nucleon recoils and correlations, and weak-magnetism effects as well as direct interactions between neutrinos of different flavors. With the given input physics (neutrino reactions and nuclear equation of state), our best simulations do not develop strong convection in the neutrino-heating layer behind the shock and do not yield explosions. With about 30\\% higher neutrino-energy deposition behind the shock, however, an explosion occurs on a timescale of 150$\\,$ms after core bounce. It leaves behind a neutron star with an initial baryonic mass of 1.4$\\,M_{\\odot}$ and ejects the $N=50$ isotopes of Sr, Y and Zr in amounts consistent with Galactic abundances. ",
        "introduction": "\\label{sec:intro}  Wilson and collaborators (\\cite[]{wilmay88}--\\cite[]{totsat98}) have routinely obtained delayed supernova explosions by the neutrino-heating mechanism~\\cite[]{betwil85} for more than ten years in one-dimensional simulations. In order to get a sufficiently large neutrino-energy deposition for reviving the stalled shock, they {\\em assumed} that neutron-finger convective instabilities boost the neutrino luminosities from the hot neutron star. Moreover, Mayle et.~al. \\cite{maytav93} used a special nuclear equation of state with a high abundance of pions in the nuclear matter to reproduce the explosion energy of SN~1987A. Both assumptions are not generally accepted. Therefore investigations of the neutrino-driven mechanism should be continued with making less controversial postulats about the involved physics.  Forced by observations of large-scale anisotropies and mixing in SN~1987A and directed by the fact that the neutrino-heating layer is convectively unstable \\cite[]{bet90}, supernova modelers began to perform multi-dimensional simulations of the post-bounce supernova evolution and discovered that violent motions set in on a timescale relevant for the delayed mechanism (\\cite[]{herben92}--\\cite[]{shiebi01}). In fact, increasing the postshock pressure by rising, outward pushing neutrino-heated matter and allowing for simultaneous accretion (thus maintaining high accretion luminosities) and shock expansion (thus enlarging the layer of neutrino-energy deposition), this postshock convection turned out to have a very helpful influence on the explosion: Models which failed to blow up in spherical symmetry were found to succeed in two dimensions.  Making use of this effect and employing an advanced version of the Herant et.~al. code \\cite{herben94}, Fryer has started to explore a variety of progenitors to determine whether neutron stars or black holes are left behind as compact remnants \\cite{fry99}. Fryer and Heger expanded the study to rotating stars \\cite{fryheg00}, and most recently Fryer and Warren have finished the first full supernova simulations in three dimensions, finding good agreement with results that were obtained in the two-dimensional case \\cite{frywar02}.  Though these simulations are undoubtedly pioneering first steps, they suffer from a major uncertainty: The transport of neutrinos in the supernova core is described in a much simplified way, using grey flux-limited diffusion and ignoring Doppler-shift and gravitational redshift effects on the neutrino spectrum. This approximation is too crude to allow for serious conclusions on the explosion mechanism. The importance of an accurate numerical description of the neutrino transport has repeatedly been pointed out by Bruenn and by Mezzacappa and collaborators (e.g., \\cite[]{mesmez98}--\\cite[]{liemes02}).  For this reason, a new generation of supernova codes has been developed in the past years with the aim to replace the widely used (multi-group) flux-limited diffusion treatments (e.g., \\cite[]{may85}--\\cite[]{coobar92}) by a direct solution of the Boltzmann transport equation (\\cite[]{yamjan99}--\\cite[]{ramjan00},\\cite[]{mesmez98,mezlie01}). Progress was also made for including the effects of general relativity (GR) in a rigorous way in neutrino-hydrodynamics codes in spherical symmetry (\\cite[]{liemez01}--\\cite[]{liemes02:code}). Although improved flux limiters for the neutrino transport \\cite[]{jan92} or two-moment closure schemes \\cite[][ and references therein]{smicer97} have been suggested and may work better than previously used flux limiting prescriptions (for an apparently successful new development in this direction, see the results by Bruenn shown in Ref.~\\cite{liemes02:code}), any such approximation has ultimately to be compared with Boltzmann solvers applied in dynamical supernova calculations.  In this conference contribution we shall summarize our efforts to combine a Boltzmann solver for the neutrino transport with a multi-dimensional hydrodynamics code, using a state-of-the-art description of neutrino-matter interactions and an approximation of GR which is (hopefully) sufficiently accurate to model supernova explosions and neutron star formation. First results of two-dimensional simulations with this new code named MuDBaTH ({\\bf Mu}lti-{\\bf D}imensional {\\bf B}oltzm{\\bf a}nn {\\bf T}ransport and {\\bf H}ydrodynamics) will also be presented. ",
        "conclusions": "\\label{sec:summary}  We have presented a set of core collapse and supernova models, which were computed in spherical symmetry and in two dimensions, using a new Boltzmann solver for the neutrino transport and improved neutrino opacities compared to the standard treatment in current supernova codes. We did not find explosions, neither without nor with postshock convection, in our most complete simulations.  Omitting the velocity-dependent terms from the neutrino momentum equation, however, leads to sufficiently large changes in the neutrino quantities and neutrino energy loss or deposition (10--30\\% between neutrinosphere and shock) that very strong convective overturn in the neutrino-heating region develops and drives a successful explosion. This demonstrates that the models are rather close to the threshold for an explosion, causing an enormous sensitivity of the post-bounce evolution to an accurate treatment of the neutrino transport and hydrodynamics. It also confirms the importance of postshock convection, because the 1D model corresponding to our exploding 2D case does not explode. When the neutrino heating is too weak, and the shock recedes to a small radius quickly after its has reached its maximum radius, however, convective activity behind the shock is suppressed and has no crucial influence on the evolution.  We also found that Ledoux convection below the neutrinosphere, although setting in shortly after bounce and persistent until the end of our simulations, occurs so deep inside the neutron star that its effects on the $\\nu_e$ and $\\bar\\nu_e$ luminosities and on the supernova dynamics are insignificant.   \\bigskip\\noindent {\\small {\\bf Acknowledgements:} We are grateful to K.~Kifonidis, T.~Plewa and E.~M\\\"uller for updates of the PROMETHEUS hydrodynamics code, to K.~Kifonidis for contributing a matrix solver suitable for parallel computer platforms, to K.~Takahashi for providing routines to calculate the improved neutrino-nucleon interactions, and to C.~Horowitz for correction formulae for the weak magnetism. We also thank M.~Liebend\\\"orfer for making output data from his simulations available to us for comparisons. The described project was supported by the Sonderforschungsbereich 375 on ``Astroparticle Physics'' of the Deutsche Forschungsgemeinschaft. The 2D simulations were only possible because a node of the new IBM ``Regatta'' supercomputer was dedicated to this project by the Rechenzentrum Garching. Computations were also performed on the NEC SX-5/3C of the Rechenzentrum Garching, and on the CRAY T90 and CRAY SV1ex of the John von Neumann Institute for Computing (NIC) in J\\\"ulich. }    \\def\\ref@jnl#1{{\\rm#1}} \\def\\apj{\\ref@jnl{ApJ}}                 % \\def\\apjl{\\ref@jnl{ApJ}}                % \\def\\apjs{\\ref@jnl{ApJS}}               % \\def\\aap{\\ref@jnl{A\\&A}}                % \\def\\aapr{\\ref@jnl{A\\&A~Rev.}}          % \\def\\aaps{\\ref@jnl{A\\&AS}}              % \\def\\physrep{\\ref@jnl{Phys.~Rep.}}   % \\def\\pra{\\ref@jnl{Phys.~Rev.~A}}        % \\def\\prb{\\ref@jnl{Phys.~Rev.~B}}        % \\def\\prc{\\ref@jnl{Phys.~Rev.~C}}        % \\def\\prd{\\ref@jnl{Phys.~Rev.~D}}        % \\def\\pre{\\ref@jnl{Phys.~Rev.~E}}        % \\def\\prl{\\ref@jnl{Phys.~Rev.~Lett.}}    % \\def\\jcomp{\\ref@jnl{J.~Comp.~Phys.}}    %"
    },
    "0212/astro-ph0212121_arXiv.txt": {
        "abstract": "I have assembled a diverse sample of galaxies from the literature with far-ultraviolet (FUV), optical, infrared (IR) and radio luminosities to explore the calibration of radio-derived and IR-derived star formation (SF) rates, and the origin of the radio-IR correlation. By comparing the 8--1000{\\micron} IR, which samples dust-reprocessed starlight, with direct stellar FUV emission, I show that the IR traces most of the SF in luminous $\\sim L_*$ galaxies but traces only a small fraction of the SF in faint $\\sim 0.01 L_*$ galaxies.  If radio emission were a perfect SF rate indicator, this effect would cause easily detectable curvature in the radio-IR correlation. Yet, the radio-IR correlation is nearly linear.  This implies that the radio flux from low-luminosity galaxies is substantially suppressed, compared to brighter galaxies.  This is naturally interpreted in terms of a decreasing efficiency of non-thermal radio emission in faint galaxies. Thus, the linearity of the \\rf correlation is a conspiracy: both indicators underestimate the SF rate at low luminosities. SF rate calibrations which take into account this effect are presented, along with estimates of the random and systematic error associated with their use. ",
        "introduction": "\\label{sec:intro}  Because ultraviolet (UV) and optical star formation (SF) rate indicators are so sensitive to dust \\citep[see, e.g.,][]{k98,adel00,uit,gold02,bell02}, there has been much recent interest in using infrared (IR) and radio luminosities in their stead \\citep[see, e.g.,][]{blain99,flores99,haarsma00,hopkins01,mann02}. While IR emission is straightforward to understand in the optically-thick case for an intensely star-forming galaxy \\citep{k98}, radio emission is a highly indirect indicator of SF rate, relying largely on the complex and poorly-understood physics of cosmic ray generation and confinement \\citep[see the excellent review by][]{condon92}. Indeed, the strongest argument for radio luminosity as a SF rate indicator has come from the astonishingly tight (a factor of two over 5 orders of magnitude in luminosity) and arguably linear radio-IR correlation \\citep[e.g.,][]{dejong85,condon91,yun01}.  This close link between the radio and IR luminosities of galaxies, even when normalized by galaxy mass \\citep[e.g.,][]{fitt88,price92}, has often been used as a supporting argument for the efficacy and robustness of radio- and IR-derived SF rates. In this paper, I compare UV, \\hans, IR and radio luminosities for a diverse sample of galaxies to demonstrate that neither the IR nor radio emissions linearly track SF rate.  I argue that the tight, nearly linear radio-IR correlation is a conspiracy: both the IR and radio luminosities of dwarf galaxies significantly underestimate the SF rate.  Finally, new SF rate calibrations which take into account this effect are presented.  \\subsection{The origin of IR and radio emission}  The primary prerequisite for an effective SF rate indicator is that it reflects the mass of young stars in some well-defined way. However, in practice, no SF rate indicator directly reflects the mass of young stars.  It is useful at this stage to develop an intuition for the physical origin, strengths and limitations of IR and radio emissions as SF rate indicators.  For more in-depth discussion of these SF rate indicators, see \\citet{k98} and \\citet{condon92}.  \\subsubsection{IR emission}  In systems with ongoing SF, the light from both newly-formed and older stars can be absorbed by dust and reprocessed into the IR. There are thus two questions that should be addressed. {\\it i)} What are the relative contributions of old and young stars to the IR luminosity?  {\\it ii)} How much light is reprocessed into the IR?  Put differently, what is the optical depth of galaxies?  Because of my focus on the young stellar populations, I will tend to focus on the optical depth of galaxies to light from young stars.  The relative balance of dust heating by young and old stars in star-forming galaxies is a matter of some debate.  One observational indicator of this balance is the temperature of the dust. Young stars in \\hii regions heat up dust to relatively high temperatures (with a low 100{\\micron} to 60{\\micron} ratio of $\\sim 1$). Older stars in the field, and far-ultraviolet (FUV) light from field OB associations (which have dispersed their natal clouds and so are relatively unattenuated in the FUV), heat the dust to much lower temperatures \\citep[100/60 $ \\ga 5$; see, e.g.,][]{lonsdale87,buat96,walterbos96}. This difference between \\hii region and diffuse dust temperatures leads to a wide range in 100{\\micron}/60{\\micron} on galaxy-wide scales, from $\\sim 10$ for quiescent early-type spiral galaxies through to $\\la 1$ for the most intensely star-forming galaxies.  This suggests that earlier types are influenced more by old stellar populations than later types; this is also supported by an analysis of far-IR (FIR) and \\ha data by \\citet{sauvage92}. For a `median' spiral galaxy, the `cold' dust IR luminosity fraction is between 50\\% and 70\\% \\citep{lonsdale87,bothun89}. Despite this domination by cooler dust, more recently it has been argued that the young, FUV-bright stars provide the dominant contribution to the IR flux \\citep[$\\sim$70\\%; see, e.g.,][]{buat96,popescu00,misiriotis01}. This is because FUV light is absorbed much more efficiently than optical light per unit dust mass.  Thus, for a `median' spiral galaxy, the IR luminosity comes from three components in roughly equal amounts: $\\sim 1/3$ of the IR is from warm dust heated by FUV light from intense SF in \\hii regions, another $\\sim 1/3$ is cold dust heated by optical photons from the old and young stellar populations, and the last $\\sim 1/3$ of the IR is from cold dust heated by FUV light from OB associations in the field \\citep{buat96}. This interesting issue is discussed further in \\S \\ref{sec:optical}.  Given the apparent dominance of young stars in determining the IR flux, it is appropriate to address the opacity of dust to light from young stars.  Observationally, there is a strong but scattered correlation between galaxy luminosity ($\\sim$ mass) and dust opacity to UV or \\ha light \\citep{wang96,adel00,uit,hopkins01,sullivan01,buat02}. Low-luminosity galaxies ($L/L_* \\sim 1/100$) tend to have substantially less dust absorption and reddening than high-luminosity galaxies ($\\sim L_*$). Furthermore, these papers demonstrate that low-luminosity galaxies have IR/FUV$\\la 1$, meaning that the IR emission of low-luminosity galaxies misses most of the SF. In contrast, many high-luminosity galaxies have IR/FUV$\\gg 1$, implying that the IR may be a relatively good SF rate indicator in this case \\citep[\\S \\ref{subsec:firfuv}]{wang96,buat02}. This will have clear implications for IR-derived SF rates, and correlations involving IR luminosities, such as the \\rf correlation.  This paper explores these implications in detail.  \\vspace{0.2cm} \\subsubsection{Radio emission} \\vspace{0.2cm}   Radio continuum emission from star-forming galaxies has two components: thermal bremsstrahlung from ionized Hydrogen in \\hii regions \\citep[see, e.g.,][]{caplan86}, and non-thermal synchrotron emission from cosmic ray electrons spiraling in the magnetic field of the galaxy \\citep[see, e.g.,][{ }for an excellent review]{condon92}. Thermal radio emission has a spectrum $\\propto \\nu^{-0.1}$, whereas non-thermal emission has a much steeper radio spectrum $\\propto \\nu^{\\alpha}$, where $\\alpha \\sim -0.8$ \\citep[however, note that $\\alpha$ can vary, and even can vary with frequency;][]{condon92}. Because of this difference in spectral shape, the relative contributions of the two emissions vary with frequency.  At lower frequencies $\\la 5$\\,GHz non-thermal radiation tends to dominate \\citep[at 1.4\\,GHz, the `standard model' of star-forming galaxies attributes typically 90\\% of the radio continuum flux of luminous spiral galaxies to non-thermal emission;][]{condon92}. Based on the standard model, thermal emission may dominate at frequencies $\\ga 10$\\,GHz \\citep[see also ][]{price92}. In addition, the relative fractions of thermal and non-thermal emission may depend on galaxy mass. Dwarf galaxies seem to have a lower non-thermal to thermal emission ratio than normal spiral galaxies \\citep{klein84,klein91,klein91b,price92}, although estimating the balance of thermal and non-thermal radio emission is painfully difficult, and can be uncertain for even well-studied galaxies at a factor of five level \\citep{condon92}. This difference between dwarf and larger galaxies is often interpreted as a higher efficiency of cosmic ray confinement in physically larger (or more massive) galaxies \\citep[e.g.,][]{klein84,chi90,price92}. For interesting discussions about the relative balance of non-thermal and thermal emission see \\citet{condon92} and \\citet{niklas97}.  \\subsection{The Radio--IR correlation}  Given the complexity of the emission mechanisms of radio continuum and IR light, it seems to be a miracle that the two fluxes are tightly correlated, with a scatter of only a factor of two. Yet, when examined closely, the \\rf correlation betrays the richness of the astrophysics which determine galaxies' radio and IR luminosities.  The slope of the \\rf correlation seems to depend on galaxy luminosity.  Samples which are richer in relatively faint galaxies ($L_{\\rm IR} \\la 10^{10} L_{\\sun}$) tend to have steep \\rf correlations in the sense that $L_{\\rm radio} \\propto L_{\\rm IR}^{\\gamma}$ and $\\gamma > 1$ \\citep[e.g.,][]{cox88,price92,xu94}, whereas samples with a better representation of highly luminous galaxies ($10^{10} L_{\\sun} \\la L_{\\rm IR} \\la 10^{12.5} L_{\\sun}$) tend to have slopes close to unity \\citep[e.g.,][]{condon91,yun01}. The differing behavior of galaxies as a function of luminosity is beautifully illustrated in Fig.\\ 5 of \\citet{yun01} and Figs.\\ 1 and 2 of \\citet{condon91}. In addition, the slope depends on the radio frequency.  At low radio frequencies $\\la 5$\\,GHz the slope tends to be steeper than unity, whereas for higher frequencies the slope approaches unity \\citep[wonderfully illustrated in Fig.\\ 2 of][]{price92}.  Workers in this field have typically sought to explain the luminosity-dependent slope in terms of heating of dust by older stellar populations, or non-thermal/thermal radio effects. \\citet{fitt88} and \\citet{devereux89} both subtracted off plausible contributions from old stellar populations (using either FIR color $\\theta$ or total IR luminosity as the constraint), which they found `linearized' the \\rf correlation. \\citet{condon91} compared IR/radio with optical $B$/radio, finding that IR-overluminous galaxies were overluminous in optical $B$-band ($\\sim 4400${\\AA}), which was interpreted as indicating contributions from old stellar populations.  \\citet{xu94} presented a model which described the non-unity slope and some of the scatter of the \\rf correlation in terms of the contributions of old stellar populations.  Similarly, a number of studies have investigated the r\\^ole of non-thermal/thermal emission on the \\rf correlation. \\citet{price92} and \\citet{nikradfir} find that thermal radio continuum (which directly reflects the SF rate) correlated linearly with the IR luminosity; however, non-thermal emission had a steeper correlation with IR luminosity with $\\gamma \\sim 1.3$. Taken together, the steepening of the \\rf correlation at low IR luminosities, and with decreasing radio frequency, have been interpreted as reflecting increasingly large contributions from old stellar population heating of the IR towards low IR luminosities, and non-thermal radio emission which is non-linearly related to the SF rate.  \\subsection{The goal of this paper}  In contrast with the commonly accepted picture, I argue that these interpretations of the \\rf correlation are incomplete because they neglect the effect of dust opacity on the IR emission of star-forming galaxies \\citep[note that][briefly discussed the role of dust opacity, but not in a luminosity-dependent sense]{lisenfeld96}. The argument can be (but has not been, as yet) pieced together from results in the literature. Empirically, high-luminosity galaxies are optically thick to FUV light, and so their IR emission reflects the SF rate reasonably well.  In contrast, low-luminosity galaxies have low IR/FUV; therefore, their IR emission underestimates the SF rate substantially \\citep{wang96}.  Yet the \\rf correlation is more or less linear \\citep[e.g.,][]{yun01}.  Therefore, the radio emission must be suppressed for low-luminosity galaxies.  This offers independent support to the argument that low-luminosity galaxies tend to have substantially suppressed non-thermal radio emission \\citep[e.g.,][]{klein84,klein91,price92}.  Thus, the \\rf correlation is linear not because both emissions reflect SF rate perfectly, but because both radio and IR emissions underestimate the SF rates of low-luminosity galaxies in coincidentally quite similar ways.  In this paper, I assemble a sample of star-forming galaxies with FUV, IR and radio data to quantitatively explore this basic argument for the first time. The result that low-luminosity galaxies have IR and radio emissions that underestimate their SF rates is not new \\citep[e.g.,][]{wang96,klein84,dale01}.  However, the assembly of an extensive star-forming galaxy sample with FUV, IR and radio data, the quantitative exploration of the consequences of this result on the \\rf correlation, and the presentation of SF rate calibrations which take into account this effect, are new.  The plan of this paper is as follows. I first investigate, in detail, dust opacity indicators, and trends in dust opacity with galaxy luminosity, in \\S \\ref{sec:sample}.  The galaxy sample is also introduced there. In \\S \\ref{sec:fuv}, the \\rf correlation is constructed, and the effect of dust opacity on the \\rf correlation is estimated. In \\S \\ref{sec:optical}, the effect of optical light from old stellar populations is discussed.  In \\S \\ref{sec:radio}, deviations from the expected trends in the radio--IR correlation are used to investigate the relationship between radio emission and SF rate.  In \\S \\ref{sec:disc}, new IR and radio SF rate calibrations are presented and discussed.  In \\S \\ref{sec:conc}, I summarize the conclusions of this study.  In Appendix \\ref{sec:app}, the FUV, optical, IR and radio data are discussed in more detail, and I present a table of galaxy photometry.  In Appendix \\ref{app:model}, I present and discuss in detail a model for a luminosity-dependent FUV optical depth. Sections 2.1--2.3, \\S 4, and the appendices are less central to my discussion of IR/radio SF rates, and may be skipped by casual readers. A distance scale compatible with $H_0 = 75$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$ is assumed, and unless stated otherwise, I correct FUV and optical data for galactic foreground extinction using \\citet{sfd}. ",
        "conclusions": "\\label{sec:conc}  I have assembled a diverse sample of galaxies from the literature with FUV, optical, IR and radio luminosities to explore the calibration of radio- and IR-derived SF rates, and the origin of the \\rf correlation.  My main conclusions are as follows.  In order to establish the efficacy of IR/FUV as an extinction indicator, I compare \\ha and 8--1000{\\micron} TIR/FUV properties of a subsample of my galaxies. I find that \\ha and FUV attenuations loosely correlate with each other, with the \\ha attenuation being roughly half of the FUV attenuation. A foreground screen model would predict an offset of a factor of a quarter.  This lends support to the claim of \\citet{calzetti94} that the nebular extinction is roughly twice that of the stellar population of the galaxy. Furthermore, when SF rates derived from TIR$+$FUV and attenuation-corrected \\ha are compared, I find that they agree to better than a factor of two (random and systematic). This strongly argues that TIR/FUV will give FUV attenuation estimates which are accurate to a factor of two, and probably much better.  Having established the efficacy of TIR/FUV as a FUV attenuation indicator, I explored trends in TIR/FUV with galaxy luminosity.  This ratio increases on average by over a factor of 30 between low-luminosity galaxies ($L \\sim 1/100 L_*$) and high-luminosity galaxies ($L \\sim 3 L_*$). Low-luminosity galaxies have TIR/FUV$ \\la 1$, meaning that they are {\\it optically thin} in the FUV. Interestingly, the gross, overall trend in TIR/FUV is naturally interpreted in terms of increasing gas surface density and galaxy metallicity with increasing galaxy mass.  Like \\citet{yun01}, I find a nearly linear \\rf correlation, with perhaps a slight tendency for faint galaxies to have a somewhat higher TIR-to-radio ratio $q_{\\rm TIR}$ than brighter galaxies. However, the strong and robust increase in TIR/FUV with luminosity would, if radio were a perfect SF rate indicator, produce a clear and easily measurable decrease in $q_{\\rm TIR}$ for fainter galaxies.  The data show the opposite (or no) trend, clearly demonstrating that {\\it radio is not a perfect SF rate indicator}. Accounting for the effects of older stellar populations using a simple FUV/optical/IR energy balance model (which is consistent with FIR color-based methods) does not change this key result.  In order to cancel out the trend in $q_{\\rm TIR}$ from optical depth effects, the non-thermal emission must be suppressed by about a factor of 2--3 in $\\sim L_*/100$ galaxies relative to $\\sim L_*$ galaxies. This result was also reached independently, using a totally different dataset and method, by \\citet{price92}. Thus, the linearity of the \\rf correlation is a conspiracy: both radio and IR underestimate the SF rate for low-luminosity galaxies.  I present SF rate calibrations which simultaneously reproduce the linearity of the \\rf correlation, and take account of the reduced non-thermal and IR emission in lower-luminosity galaxies. However, there is considerable scatter in the SF rate calibrations, which can exceed a factor two at low galaxy luminosities.  This highlights the possible influence of selection effects in interpreting the IR or radio emission from distant galaxies. Another challenge for those wishing to estimate the SF rates of distant galaxies is the non-trivial physics that links IR/radio and SF rate.  For example, the evolution of dust opacity, the importance of old stellar populations, and the evolution of the efficiency of cosmic ray confinement are all essentially unconstrained as a function of lookback time. This adds considerable systematic uncertainty to our understanding of galactic SF rates in the distant Universe."
    },
    "0212/astro-ph0212298_arXiv.txt": {
        "abstract": "We have examined the complete set of X-ray afterglow observations of dark and optically bright GRBs performed by BeppoSAX until February 2001. X-ray afterglows are detected in $\\sim 90\\%$ of the cases. We do not find significant differences in the X-ray spectral shape, in particular no higher X-ray absorption in GRBs without optical transient ( dark GRBs ) compared to GRBs with optical transient ( OTGRBs ). Rather, we find that the 1.6-10 keV flux of OTGRBs is on average about 5 times larger than that of the dark GRBs. A K-S test shows that this difference is significant at 99.8\\% probability. Under the assumption that dark and OTGRB have similar spectra, this could suggest that the first are uncaught in the optical band because they are just faint sources. In order to test this hypothesis, we have determined the optical-to-X ray flux ratios of the sample. OTGRBs show a remarkably narrow distribution of flux ratios, which corresponds to an average optical-to-x spectral index $\\overline{\\alpha}_{OX} ^{ OT}= 0.794\\pm 0.054 $. We find that, while 75\\% of dark GRBs have flux ratio upper limits still consistent with those of OT GRBs, the remaining 25\\% are 4 - 10 times weaker in optical than in X-rays. The significance of this result is $\\geq 2.6\\sigma$. If this sub-population of dark GRBs were constituted by objects assimilable to OTGRBs, they should have shown optical fluxes higher than upper limits actually found. We discuss the possible causes of their behaviour, including a possible occurrence in high density clouds or origin at very high redshift and a connection with ancient, Population III stars. ",
        "introduction": "About 50\\% of well-localized GRBs show optical transients (OTs) successive to the prompt gamma-ray emission, whereas an X-ray counterpart is present in 90\\% of cases.  It is possible that late and shallow observations could not detect the OTs in some cases; several authors argue that dim and/or rapid decaying transients could bias the determination of the fraction of truly obscure GRBs \\citep{Fyn01a,Ber02}. However, recent reanalysis of optical observations \\citep{Rei01,Ghi00,Laz00} has shown that GRBs without OT detection (usually dark GRBs, FOAs Failed Optical Afterglows, or GHOSTs, Gamma ray burst Hiding an Optical Source Transient) have had on average weaker optical counterparts, at least 2 magnitudes in the R band, than GRBs with OTs. Therefore, they appear to constitute a different class of objects, albeit there could be a fraction undetected for bad imaging.  Two hypothesis have been put forward to explain the behaviour of GHOSTs. First, they are similar to the other bright GRBs, except for the fact that their lines of sight pass through large and dusty molecular clouds, that cause high absorption. Second, they are more distant than GRBs with OT, at $ z \\ga 5 $ \\citep{Fru99}, so that the Lyman break is redshifted into the optical band. These GRBs might be associated with the explosion of ancient Population III, high mass stars. Nevertheless, the distances of a few dark GRBs have been determined and they do not imply high redshifts \\citep{Djo02,Ant00,Pir02}  Goal of this paper is an analysis of a complete sample of BeppoSAX X-ray afterglows in order to distinguish between these various scenarios, including all x-ray fast observation from the launch to February 2001. In \\S 2 and \\S 3 we present the data analysis of the afterglows and we show the results, whose implications are discussed in \\S 4. Finally, we summarize our conclusions in \\S 5. ",
        "conclusions": "We have discussed the issue of GRBs with X-ray but no optical afterglows. We have performed a standard temporal and spectral analysis of a complete sample of 31 GRB X-ray follow up observations of BeppoSAX, i.e. all the fast observations from the launch until February 2001. We have found that X-ray afterglows follow the prompt gamma emission in $84\\% - 94\\% $ of the cases.  We have obtained the 1.6 - 10 keV fluxes 11 hours after the trigger for each GRB and the values of N$_{H}$ at $z=1$ to compare the absorption properties for strong X-ray afterglows. While absorption of optically bright and dark GRBs does not appear to be significatively different, the fluxes of GRBs with OT are on average about 5 times stronger than GHOST ones. The probability that GHOSTs and optically bright GRBs belong to the same population in fluxes is $\\leqslant 0.002$.  From the very fact that X-ray flux of dark GRBs is 5 times lower than OTGRBs, the optical flux could be $\\sim 2$ magnitude lower under the assumption that the shape of the optical-to-X spectrum is the same of OTGRBs. This difference could explain the non-detection of the optical transient. In order to test this hypothesis, we have calculated the optical-to-X flux ratios of OTGRBs and upper limits for GHOSTs. OTGRBs show a tight correlation of optical and X-ray fluxes. The mean for OTGRBs is $<$ log $f_{OX} ^{ OT}$ $>^{m}$ = $0.3\\pm 0.22$ and $\\sigma = 0.4 $ ; the probability of a chance correlation is a marginal $\\sim 1\\%$. We find that $75\\%$ of GHOSTs have $f_{OX}$ upper limits similar to OTGRB ones; however, the remaining $\\sim 25\\%$ of dark bursts are fainter in optical than in X-rays, being their average optical-to-x flux ratio $<$ log $f_{OX} $ $>^{m}$ $\\leq -0.4$. Thus, we have a strong indication that for these bursts the spectra are different from OTGRBs. This result is significant at $\\geq 2.6 \\sigma$ level.  Two different interpretations for this effect can be given: 1) location at $z>5$ ; 2) higher absorption. In the very high redshift scenario, the optical flux of the sample is extinguished by the intervening Ly-$\\alpha$ systems, while the X-ray flux lower than OTGRBs is understood in terms of a higher distance.  However, given the fact that some GHOSTs of the sample almost certainly do not lie at very high redshift, we have considered the alternative possibility of occurrence in dusty and dense environments like GMCs. We have not found that these bursts have a higher absorption than optically bright GRBs, but we note that upper limits on $N_{H}$ are consistent with those of giant clouds. In the case of GRB000210 our model-independent analysis has shown a depletion in the optical which is compatible with the X-ray absorption measured by Chandra, assuming a dust-to-gas ratio similar to that of our Galaxy.  In the near future, a key role will be played by fast and deep follow up X-ray and optical observations of GRBs, which will allow us to constrain better their spectral properties. In particular, observations in the IR band are a really important tool because they are less sensitive to dust and to Ly-$\\alpha$ extinction. They will enable us to investigate dark GRB properties like distance, that is a crucial piece of information to disclose the nature of these objects."
    },
    "0212/astro-ph0212065.txt": {
        "abstract": "With most physicists and astrophysicists in agreement that black holes do indeed exist, the focus of astrophysical black hole research has shifted to the detailed properties of these systems. Nature has provided us with an extremely useful probe of the region very close to an accreting black hole --- X-ray irradiation of relatively cold material in the vicinity of the black hole can imprint characteristic features into the X-ray spectra of black hole systems, most notably the K$\\alpha$ fluorescent line of iron. Detailed X-ray spectroscopy of these features can be used to study Doppler and gravitational redshifts, thereby providing key information on the location and kinematics of the cold material. This is a powerful tool that allows us to probe within a few gravitational radii, or less, of the event horizon.  Here, we present a comprehensive review of relativistic iron line studies for both accreting stellar mass black holes (i.e., Galactic Black Hole Candidate systems; GBHCs), and accreting supermassive black holes (i.e., active galactic nuclei; AGN).  We begin with a pedagogical introduction to astrophysical black holes, GBHCs, AGN, and accretion disks (including a brief discussion of recent work on the magnetohydrodynamical properties of accretion disks).  We then discuss studies of relativistic iron lines in the AGN context, and show how differences between classes of AGN can be diagnosed using X-ray spectroscopy.  Furthermore, through a detailed discussion of one particular object (MCG--6-30-15), we illustrate how the exotic physics of black hole spin, such as the Penrose and Blandford-Znajek processes, are now open to observational study.  We proceed to discuss GBHCs, which turn out to possess rather more complicated X-ray spectra, making robust conclusions more difficult to draw. However, even in these cases, modern X-ray observatories are now providing convincing evidence for relativistic effects.  We conclude by discussing the science that can be addressed by future X-ray observatories. ",
        "introduction": "\\label{Intro}  Gravitational collapse of normal matter can produce some of the most exotic objects in the universe --- neutron stars and black holes. Proving that these objects exist in Nature occupied theoretical and observational astrophysicists for much of the 20th century.  Most of the detailed debate centered around understanding the possible final states of massive stars.  On his now famous sea voyage from India to England in 1930, Subrahmanyan Chandrasekhar considered the structure of white dwarf stars --- compact stellar remnants in which gravitational forces are balanced by electron degeneracy pressure.  He realized that, if the white dwarf was sufficiently massive, the degenerate electrons will become relativistic thereby rendering the star susceptible to further gravitational collapse \\cite{chandrasekhar:31a,chandrasekhar:31b}.  Although hotly debated by Arthur Eddington, Chandrasekhar correctly deduced that a white dwarf would undergo gravitational collapse if its mass exceeded $M_{\\rm ch}\\approx 1.4\\Msun$ (where $1\\Msun=2.0\\times 10^{33}\\g$ is the mass of the Sun), a limit now known as the Chandrasekhar limit\\footnote{The Russian physicist Lev Landau independently calculated the upper mass limit of a white dwarf star, and his results were published in 1932 \\cite{landau:32a}.}.  Once gravity overwhelms electron degeneracy pressure, it is thought that neutron degeneracy pressure is the last, best hope for averting total gravitational collapse.  Objects in which gravitational forces are balanced by neutron degeneracy pressure are called {\\it neutron stars}.  Although there was initial hope that nuclear forces would always be sufficient to resist gravity, the upper limit to the mass of a neutron star is now believed to be in the range $M\\approx 1.8-2.2\\Msun$\\cite{akmal:98a,srinivasan:02a}.  Uncertainties arising from the equation of state at super-nuclear densities continue to plague our determination of this critical mass, but an absolute upper limit of $M\\approx 4\\Msun$ arises from very general considerations, i.e., the validity of General Relativity and the principle of causality\\cite{rhoades:74a}.  Above this mass, it is thought that complete gravitational collapse cannot be avoided.  In particular, Hawking's singularity theorems\\cite{hawking:73a} show that the formation of a spacetime singularity is unavoidable (irrespective of the mass/energy distribution) once the object is contained within the light trapping surface.  The result is a {\\it black hole}, i.e., a region of spacetime bounded by an event horizon and, at its heart, possessing a spacetime singularity.  While the above considerations now have a firm theoretical base, observational astrophysics was, and continues to be, critically important in guiding our understanding of such extreme objects.  In the case of both neutron stars and black holes, the very existence of these objects was only widely accepted when compelling observational evidence was forthcoming.  For neutron stars, the pivotal observation was the discovery of pulsars by Jocelyn Bell and Anthony Hewish via radio observations taken from Cambridge.  Black holes gained wide acceptance after it was demonstrated that the X-ray emitting compact object in the binary star system Cygnus~X-1 did, in fact, possess a mass in excess of the maximum possible neutron star mass\\cite{murdin:71a,hutchings:78a,balog:81a,gies:86a,herrero:95a}. This made it the first of the so-called Galactic Black Hole Candidates (GBHCs), a class that has now grown to include some two dozen objects.  We now know of another class of black holes --- the supermassive black holes, with masses in the range of $10^5-10^{9.5}\\Msun$, that reside at the dynamical centers of most, if not all, galaxies\\footnote{There have been recent suggestions of a possible population of ``intermediate mass'' black holes, with masses in the range of $10^3$--$10^5$\\Msun.  (See the discussion in \\cite{miller:02d}, for example.)  The existence of such objects is still speculative, and detailed studies of their spectra, of the kind to be discussed in this review, do not yet exist.  We shall therefore not consider these objects further in this review.}.  Today, by far the strongest case for a supermassive black hole can be made for our own Galaxy. Modern high-resolution, infra-red imaging reveals that the stars in the central-most regions of our Galaxy are orbiting an unseen mass of $2.6\\times 10^6\\Msun$\\cite{eckart:97a,ghez:98a,schoedel:02a}. Furthermore, studies of the orbital dynamics (which now include measured accelerations as well as velocities; \\cite{ghez:00a,eckart:02a}) constrain the central mass to be extremely compact.  According to conventional physics, the only long-lived object with these properties is a supermassive black hole. Alternatives, such as a compact cluster of neutron stars, would suffer a dynamical collapse on short time scales \\cite{genzel:97a}.  \\begin{figure}[t] \\label{fig:tanaka} \\centerline{ \\hspace{2cm} \\psfig{figure=tanaka.ps,angle=270,width=1.0\\textwidth} } \\caption{Continuum subtracted iron line from the long July-1994 {\\it ASCA} observation of the Seyfert-1 galaxy MCG--6-30-15 \\cite{tanaka:95b}. The dashed line shows a model consisting of iron line emission from a relativistic accretion disk around a non-rotating (Schwarzschild) black hole, with a disk inclination of $i=30^\\circ$, and an emissivity profile of $r^{-3}$ extending down to the radius of marginal stability ($6\\,{\\rm GM/c^2}$).} \\end{figure}   Having established beyond reasonable doubt that black holes exist, it is obviously interesting to perform detailed observational studies of them.  The regions in the immediate vicinity of a black hole bear witness to complex interactions between matter moving at relativistic velocities, electromagnetic fields, and the black hole spacetime itself.  Given that the apparent angular scales of even the biggest black hole event horizons are $\\sim 10^{-6}\\arcsec$, direct imaging studies of these regions will not be possible for many years\\footnote{However, on time scales of 10--20 years, there are plans for both radio wavelength and X-ray interferometers which will be able to image nearby supermassive black holes on spatial scales comparable with that of the event horizon (see \\S\\ref{sec:maxim}).}.  In the meantime, we must study these regions using more indirect methods, chief among which are spectroscopic methods.  As we will detail in this review, Nature has provided us with a well-understood and extremely useful spectral diagnostic of matter in the near vicinity of astrophysical black holes.  In essence, relatively cold matter in the near vicinity of an astrophysical black hole will inevitably find itself irradiated by a spectrum of hard X-rays \\cite{guilbert:88a,lightman:88a}.  The result can be a spectrum of fluorescent emission lines, the most prominent being the K$\\alpha$ line of iron at an energy of $6.40-6.97\\keV$ (depending upon the ionization state of the iron) \\cite{basko:78a,george:91a,matt:91a}. Ever since the launch of the {\\it Advanced Satellite for Cosmology and Astrophysics} (ASCA) in February 1993, X-ray astrophysicists have had the capability to identify this emission line and measure its spectral profile.  Figure~\\ref{fig:tanaka} shows the iron line in the X-ray emissions originating near the supermassive black hole in the galaxy MCG--6-30-15 \\cite{tanaka:95b}.  Bearing in mind that the line is intrinsically narrow with a rest-frame energy of 6.4\\,keV, it can be seen that the line has been dramatically broadened and skewed to low-energies.  It is now widely accepted that the line originates from material that is just a few gravitational radii from the black hole, and possesses a profile that is shaped by (relativistic) Doppler shifts and gravitational redshift effects.  Investigating these spectral features in X-ray luminous black hole systems has given us the clearest window to date on the physics that occurs in the immediate vicinity of astrophysical black holes.  The intent of this review is to describe our current understanding of black hole astrophysics, with an emphasis on what has been learnt by utilizing these X-ray spectral signatures.  We begin by discussing the basic theoretical framework within which we understand the astrophysical environment around black holes.  A central and important part of this discussion is an introduction to the modern theory of accretion disks.  Hand-in-hand with the theoretical discussion, we will introduce the necessary phenomenology associated with stellar mass and supermassive black holes.  We then describe the array of past, current, and future X-ray observatories which have bearing on relativistic studies of black holes before discussing how iron line spectroscopy has dramatically improved our current understanding of black hole astrophysics.  We conclude by presenting the prospects for future research in this field. ",
        "conclusions": "The past few years have seen a shift in the mind-set of many observationally-oriented black hole researchers.  A small number of important observations (the mass of the compact object in the binary star system Cyg~X-1, stellar motions in our Galactic Center, the gas disk in M87, and the maser disk in NGC~4258) have led most astronomers to the conclusions that black holes do indeed exist beyond any reasonable doubt.  With this established, the interest has shifted to the demographics of black holes in the Universe, and the detailed astrophysics of black hole systems.  As we have described, X-ray spectroscopy currently provides the best understood method of exploring the astrophysical environment in the immediate vicinity of an accreting black hole.  The accretion disk is the engine that drives all observable phenomena from accreting black hole systems.  In many systems, the surface layers of the accretion disk are expected to produce X-ray spectral features (so-called X-ray reflection signatures) in response to external hard X-ray illumination from a disk-corona --- the most prominent spectral feature is often the fluorescent K$\\alpha$ emission line of iron.  This line, which has an intrinsically small energy width, will be dramatically broadened and skewed by the rapid orbital motions of the accreting material and strong gravity of the black hole.  These relativistic spectral features can be identified {\\it unambiguously} in many of the AGN for which X-ray data of sufficient quality exist.  Similar spectral features can also be found in many AGN and GBHCs for which poorer quality data exist --- while it may not be possible to rule out other forms of spectral complexity in these objects, relativistic disk features are often the most compelling and physical of the possibilities.  Given how generic these features are, they provide a powerful way to study the near environment of accreting black holes.  Moderate luminosity radio-quiet AGN, the Seyfert galaxies, present the cleanest examples of these spectral features.  Analysis of X-ray data from {\\it ASCA} and, more recently, {\\it XMM-Newton} finds evidence for a rather cold accretion disk extending all of the way down to the radius of marginal stability around the black hole.  In one case (MCG--6-30-15), the data argue strongly that the black hole is rapidly-rotating.  Furthermore, there may be suggestions that the accretion disk is torqued by processes associated with the spinning black hole and, in particular, may be tapping into the rotational energy of the spinning black hole. The case of MCG--6-30-15, which we discussed in some detail, illustrates the exotic black hole physics that can be addressed via these techniques.  The difference in iron line properties between Seyfert galaxies and other types of AGN allows us to probe the physical state of the accretion flow as a function of, in particular, the mass accretion rate.  The fact that higher-luminosity AGN display weaker spectral signatures is very likely due to increasing ionization of the disk surface as the accretion rate (relative to that rate which would produce the Eddington luminosity) is increased.  The weakness of the X-ray reflection displayed by radio-loud AGN and low-luminosity AGN is much less certain (primarily due to the paucity of high quality datasets).  Possibilities include ionization of the disk surface, dilution of the disk spectrum by X-ray emission from a relativistic jet, and/or the transition of the disk into a radiatively-inefficient state (which would be extremely hot and optically-thin thereby producing no X-ray reflection features).  Although it is somewhat counter-intuitive given that they can be three orders of magnitude brighter than AGN, it has been much more difficult to study relativistic disk signatures in GBHCs.  In addition to the fact that the X-ray spectra of GBHC are inherently more complex (since the X-ray band contains the thermal disk emission and the disk is almost always strongly ionized), the high X-ray fluxes from these sources can readily overwhelm sensitive photon counting spectrometers. However, with {\\it Chandra} and {\\it XMM-Newton} we can now study relativistic X-ray reflection in GBHCs free of these instrumental constraints, at medium-to-high spectral resolution, and high signal-to-noise.  Consequently, relativistic iron lines have been found in several GBHCs with a variety of ``spectral states''.  With further study of these features over the coming few years, we should be able to constrain the gross geometry of the accretion flow as a function of spectral state.  The future is bright.  Until just a few years ago, the physics of relativistic accretion disks and the astrophysics of black hole spin was firmly in the realm of pure theory.  Modern X-ray spectroscopy has given us a powerful tool with which we can peer into this exotic world.  There is still much to be learnt --- much of the physics determining the properties of accretion disks, the demographics of black hole spin (which is closely related to black hole formation and history), and the prevalence of spin-energy extraction (which is of fundamental interest) are just a few of the open questions.  Continued investigation with increasingly capable instruments promise to answer these, and many more, questions."
    },
    "0212/astro-ph0212537_arXiv.txt": {
        "abstract": "We construct an evolutionary spectral energy distribution (SED) model of a starburst region, from the ultraviolet to submillimetre wavelengths. This model allows us to derive the star formation rate, optical depth by dust and apparent effective radius of starburst regions at various wavelengths; as a result, the intrinsic surface brightness of starburst regions can be derived. Using this SED model, we analyse 16 UV-selected starburst galaxies and 10 ultraluminous infrared galaxies. The derived star formation rates and optical depths are compared with emission line measurements and found to be consistent. The derived apparent effective radii are also consistent with observations. From the SED analysis, we find a bimodal property of the star formation rate with the optical depth and the compactness of stellar distributions. While mild starbursts have a limiting intrinsic surface brightness $ L_{bol} r_e^{-2} \\simeq 10^{12}$ L$_{\\odot} $kpc$^{-2}$, intense starbursts tend to be more heavily obscured and concentrated within a characteristic scale of $ r_e \\simeq 0.3 $ kpc. We suggest that the mild starbursts can be triggered by a self-gravitating disc instability in which feedback is effective in the shallow gravitational potential. On the other hand, the intense starbursts can be induced via an external dynamical perturbation like galaxy merging, in which feedback is less effective due to the deep gravitational potential attained by the large gas concentration within the central starburst region. ",
        "introduction": "Starbursts are an intensive mode of star formation in galaxies (Storchi-Bergmann, Calzetti \\& Kinney 1994; Kennicutt 1998). Ultraluminous infrared galaxies (ULIRGs) with $L_{\\rm IR} \\ge 10^{12}$ L$_{\\odot}$ are probably the most active and luminous starburst galaxies. Indeed recent ISO observations suggest that many ULIRGs are powered by starbursts with little contribution from AGN (Genzel et al. 1998; Lutz et al. 1998).  On the other hand, starburst galaxies of lower activity are HII galaxies and blue compact dwarf galaxies, known as UV-selected starburst galaxies (UVSBGs; Kinney et al. 1993; Storch-Bergmann et al. 1994; Calzetti, Kinney \\& Storchi-Bergmann 1994; McQuade, Calzetti \\& Kinney 1995; Gordon, Calzetti \\& Witt 1997; Takagi, Arimoto \\& Vansevi\\v cius 1999).   Recently, an upper limit of `bolometric surface brightness' was given by Meurer et al. (1997) for a sample of starburst galaxies observed in the rest frame 1) UV, 2) far infrared (FIR) and H$\\alpha$, and 3) 21 cm radio continuum emission. This limit seems to be physically associated with an instability in gaseous discs (e.g. Toomre 1964; Quirk 1972) as Kennicutt (1989) had already suggested for the star formation in normal disc galaxies. In their study, the effective radii are determined from observations at the different wavelengths for each galaxy. However,  if a galaxy is optically thick the observed effective radius strongly depends on the degree of extinction, since the observed light comes out solely from the region in which the optical depth is less than unity. Moreover, many ULIRGs have multiple starburst regions (e.g. Scoville et al. 2000; Surace \\& Sanders 2000), while most of UVSBGs show only one major central starburst region. These effects suggest that the panchromatic intensity limit on starbursts proposed by Meurer et al. (1997) may not be appropriate for ULIRGs.   To determine the intensity limit of starbursts with various levels of activity ranging from UVSBGs to ULIRGs, it is necessary to estimate both star formation rates (SFRs) and effective radii in a unified manner, whatever the degree of extinction. While typical $V$-band optical depths in UVSBGs are $\\tau_V \\sim 0.3 - 2$ (Takagi, Arimoto, \\& Vansevi\\v cius 1999; Storchi-Bergmann et al. 1994), those in ULIRGs reach $\\tau_V \\sim 5-50$ (Genzel et al. 1998). Such a wide variation in $\\tau_V$ causes a serious problem in deriving SFRs, since none of the currently used SFR indicators, such as H$\\alpha$ (or other Balmer lines), UV continuum, or FIR luminosities, are commonly applicable for starbursts ranging from UVSBGs to ULIRGs. The FIR luminosity is not a good measure of bolometric luminosity when starburst regions are optically thin, while H$\\alpha$ luminosity is difficult to use if the mean surface gas density becomes larger than 50 M$_\\odot$ pc$^{-2}$, corresponding to $\\tau_V \\sim 2$, causing a significant extinction at H$\\alpha$ (Kennicutt 1998). What is worse, by using Br$\\gamma$ and FIR luminosities, Kennicutt (1998) has shown that the FIR luminosities give systematically higher SFRs (by a factor of $\\sim 2$) than the Br$\\gamma$s. Sullivan et al. (2000) compared the H$\\alpha$- and UV-derived SFRs of nearby starbursts to find that starburst galaxies are typically over-luminous in the UV for a given H$\\alpha$ luminosity. The effect is strongest in the less luminous galaxies.  When comparing ULIRGs and UVSBGs, the bolometric surface brightness should be derived for each individual starburst region in ULIRGs. The wealth of ground-based low-resolution imaging surveys have provided little information on the stars in the starburst region near the confusion limit $ < 1''$.  Even the {\\em Hubble Space Telescope} ({\\it HST}) with its higher resolution $ <0.1''$ fails to go deep enough into the starburst regions, which are obscured by dust at optical wavelengths.  The Near-Infrared Camera and Multi-Object Spectrometer (NICMOS) on {\\it HST} can observe the obscured starburst regions more directly at near infrared (NIR) wavelengths where the effects of dust extinction are reduced significantly compared to visual wavelengths. Scoville et al. (2000) studied the morphologies of 24 ULIRGs with NICMOS and found that light profiles of nine of them were well fit by an $r^{1/4}$ law, rather than an exponential profile. The apparent effective radii of ULIRGs tend to be more compact than the extent of gas observed with high-resolution imaging of the CO emission (Bryant $\\&$ Scoville 1999).  This suggests that the distribution of stars is more concentrated than that of dust and gas in the starburst regions. Thus, the dimming of light due to the surrounding dust can be so large that the real distribution of stars deviates from the apparent distribution of the light even the longest observable wavelength with NICMOS of 2.2 $\\mu$m.    It is thus crucial to establish a measure of both SFRs and the spatial distribution of stars which can be applied to starbursts of any optical depth. Clearly, a new recipe to derive both SFRs and the stellar distributions is required to investigate the physical properties of starbursts comprehensively.  In this paper, we will use spectral energy distributions (SEDs) from the far ultraviolet (FUV) to submillimetre (submm) to derive SFRs of starburst regions. The SED from the UV to submm can represent the bolometric luminosity, irrespective of the dust extinction. We construct an `evolutionary' SED model for starburst regions from a simple but realistic point of view.  We show that our SED model can explain a wide variety of SEDs from optically thin starbursts (UVSBGs) to optically thick ones (ULIRGs). It is important to note that our SED model can derive not only the SFR, but also the other starburst properties like the optical depth. Moreover, we can derive the apparent effective radius of starbursts at various wavelengths for a given geometry, since we properly take into account the radiative transfer in a dusty medium. Therefore, the evolutionary stages of starburst galaxies can be investigated in a unified manner for a whole range of activity, which can then be compared with the observations of the SED and the effective radius, directly. The systematic estimation of the SFRs and the effective radius can give an answer to the question: is the panchromatic limit of starburst intensity reported by Meurer et al. (1997) also valid for ULIRGs?   This paper is organized as follows. In Section 2, we describe our evolutionary SED model for starburst galaxies. In Section 3, we summarize the model properties.  In Section 4, we apply our model to a sample of UVSBGs and ULIRGs in the local Universe and describe the resulting properties of nearby starburst galaxies derived from our SED model fitting. Sections 5 and 6 give discussions and conclusions, respectively. ",
        "conclusions": ""
    },
    "0212/gr-qc0212093_arXiv.txt": {
        "abstract": "By using the complex angular momentum approach, we prove that the quasinormal mode complex frequencies of the Schwarzschild black hole are Breit-Wigner type resonances generated by a family of surface waves propagating close to the unstable circular photon (graviton) orbit at $r=3M$. Furthermore, because each surface wave is associated with a given Regge pole of the $S$-matrix, we can construct semiclassically the spectrum of the quasinormal-mode complex frequencies from Regge trajectories. The notion of surface wave orbiting around black holes thus appears as a fundamental concept which could be profitably introduced in various areas of black hole physics in connection with the complex angular momentum approach. ",
        "introduction": "Since the pioneering work of Watson \\cite{Watson18} dealing with the propagation and diffraction of radio waves around the Earth, the complex angular momentum (CAM) method has been extensively used in several domains of scattering theory (see the monographs of Newton \\cite{New82} and of Nussenzveig \\cite{Nus92} and references therein for various applications in quantum mechanics, nuclear physics, electromagnetism, optics, acoustics and seismology). The success of the CAM method is due to its ability to provide a clear description of a given scattering problem by extracting the physical information (linked to the geometrical and diffractive aspects of the scattering process) which is hidden in partial-wave representations.  The CAM method was first used in gravitational physics by Chandrasekar and Ferrari \\cite{Chandra} in their study of nonradial oscillations of relativistic stars. They used the theory of Regge poles \\cite{New82} to determine the flow of gravitational energy through the star. The general framework for the CAM description of Schwarzschild black hole scattering was developed by Andersson and Thylwe \\cite{Andersson1}, which Andersson \\cite{Andersson2} then used to interpret the black hole glory. An important concept, which is naturally present in the CAM point of view, is that of a surface wave, which allows a description of the diffractive effects of scattering. Andersson established that the surface waves orbiting around the Schwarzschild black hole of mass $M$ propagate close to the unstable photon orbit at $r=3M$. Yet except for these articles, the CAM/surface-wave approach to resonant scattering in black hole physics has been neglected in favor of the quasi-normal modes (QNMs), in which the dynamical response to an external perturbation is explained in terms of resonant frequencies. For recent reviews of the QNM approach see Refs.~\\onlinecite{Kokkotas} and~\\onlinecite{Nollert} as well as chapter 4 of Ref.~\\onlinecite{Frolov-Novikov}. An introduction to black hole scattering can be found in \\cite{AndJens}.    In this paper we will establish the connection between Andersson's surface waves and the QNMs. More precisely, we shall prove here that the QNM complex frequencies of the Schwarzschild black hole are Breit-Wigner-type resonances generated by the surface waves. Moreover, because each surface wave is associated with a given Regge pole of the $S$-matrix, we can construct the spectrum of the QNM complex frequencies from the Regge trajectories, i.e., from the curves traced out in the CAM plane by the Regge poles as a function of the frequency.  As early as 1972, it was suggested by Goebel \\cite{Goebel} that the black hole normal modes could be interpreted in terms of gravitational waves in spiral orbits close to the unstable photon orbit at $r=3M$, which decay by radiating away energy. In the present paper, using the CAM approach, we establish this appealing and physically intuitive picture on a rigorous basis. To conclude, we also provide a framework for future developments. ",
        "conclusions": "To conclude, we would like first to comment on some aspects of our work and then to consider some possible extensions of the CAM approach in gravitational physics:  - We have established the connection between Andersson's surface waves and the QNM complex frequencies of the Schwarzschild black hole in the particular context of wave scattering. This connection is more general and could be also obtained directly from the Regge-Wheeler equation, by extending the CAM method developed by Sommerfeld \\cite{Sommerfeld49} as an alternative to the Watson approach \\cite{Watson18}.  - Because the gravitational radiation created in many black hole processes is dominated at intermediate timescales by QNMs, it can be always interpreted in terms of surface waves. Whatever the perturbation which modifies the geometry of a Schwarzschild black hole, it leads to the excitation of surface waves localized close to the unstable photon orbit. During its repeated circumnavigations, a given surface wave decays by radiating away its energy, leading to the damped ringing of the geometry. This mechanism is analogous to the damped ringing of the Earth which occurs during several days after a large earthquake and which can be explained in term of the Rayleigh surface wave. \\begin{figure} \\includegraphics[height=6cm,width=8cm]{fig_3} \\caption{\\label{fig:suivi-s2}The Regge poles $\\lambda_n(\\omega)$ followed for $\\omega=0\\to 5$, $n=1,2,3,4$, for gravitational perturbations ($2M=1$). } \\end{figure} \\begin{table} \\caption{\\label{tab:table2} A sample of QNM frequencies for gravitational perturbations ($2M=1$).} \\begin{ruledtabular} \\begin{tabular}{cccccc} &       & \\quad Exact\\footnotemark[1] \\quad & \\quad Exact\\footnotemark[1] \\quad &   Semiclassical  &  Semiclassical   \\\\ $\\ell $ & $n$ &$   \\omega ^{(o)}_{\\ell n}  $ &  $  \\Gamma _{\\ell n}/2  $ & $  \\omega ^{(o)}_{\\ell n}   $ &  $ \\Gamma _{\\ell n}/2  $  \\\\ \\hline 2& 1 & 0.74734  &-0.17793 & 0.75812 &  -0.17644  \\\\ & 2 & 0.69342  &-0.54783 & 0.78134 &  -0.49976  \\\\ & 3 & 0.60211  &-0.95655 & 0.77683 &  -0.83126   \\\\ 3& 1 & 1.19889  &-0.18541 & 1.20332 &  -0.18455  \\\\ & 2 & 1.16529  &-0.56260 & 1.20089 &  -0.54409   \\\\ & 3 & 1.10337  &-0.95819 & 1.18284 &  -0.89846  \\\\ 4& 1 & 1.61836  &-0.18832 & 1.62052 &  -0.18792  \\\\ & 2 & 1.59326  &-0.56866 & 1.61119 &  -0.56020  \\\\ & 3 & 1.54542  &-0.95982 & 1.58849 &  -0.92907 \\end{tabular} \\end{ruledtabular} \\footnotetext[1]{from Chapter 4 of Ref.~\\onlinecite{Frolov-Novikov} or Refs.~\\onlinecite{Kokkotas} and~\\onlinecite{Nollert}.} \\end{table}    - The CAM method could be naturally introduced in many other areas of the physics of relativistic stars and black holes. Such a program could include in particular the following topics:  \\noindent (i) The interpretation of the QNM of neutron stars and of the Kerr black hole in terms of surface waves. For the Kerr black hole, such an approach could easily explain the splitting of the quasinormal frequencies due to rotation.  \\noindent (ii) The analysis of Hawking radiation as well as of Kerr black hole superradiance.  \\noindent (iii) The study of black holes immersed in asymptotically anti-de Sitter space-times. The CAM approach could provide new tests of the AdS/CFT correspondence recently proposed in the context of superstring theory \\cite{Maldacena}. With this aim in view, the (2+1)-dimensional BTZ black hole \\cite{btz} seems to us very interesting because, in that particular space-time, the wave equation can be solved exactly (see, e.g., \\cite{Cardoso}) and therefore Regge poles can be analytically obtained.   \\noindent (iv) The study of artificial black holes \\cite{Novello}. Here, it would be possible to benefit from the formidable CAM machinery developed in electromagnetism, optics and acoustics."
    },
    "0212/astro-ph0212471_arXiv.txt": {
        "abstract": "Quintessence models leading to a constant equation of state are studied in hyperbolic universes. General properties of the quintessence potentials $V(\\phi)$ are discussed, and for some special cases also the exact analytic expressions for these potentials are derived. It is shown that the observed angular power spectrum of the cosmic microwave background (CMB) is in excellent agreement with some of the quintessence models even in cases with negative curvature. It is emphasized that due to a $(w_\\phi, \\Omega_\\phi, \\Omega_{\\hbox{\\scriptsize c}})$-degeneracy a universe with negative spatial curvature cannot be excluded. ",
        "introduction": "The recent observations of the anisotropy of the cosmic microwave background (CMB) \\cite{Netterfield_et_al_2001,Lee_et_al_2001,Halverson_et_al_2001} together with the power spectrum of the large scale structure (LSS) \\cite{Peacock_Dodds_1994,Hamilton_Tegmark_Padmanabhan_2000,% Percival_et_al_2001,Efstathiou_2dF_team_2002}, and the magnitude-redshift relation of the supernovae Ia \\cite{Hamuy_et_al_1996,Riess_et_al_1998,Perlmutter_et_al_1999} give strong evidence that the present mean energy density $\\varepsilon_{\\hbox{\\scriptsize tot}}$ of the universe consists not only of radiation, baryonic and cold dark matter, but also of a dominant component with negative pressure which nowadays is called dark energy. An obvious candidate for this new energy is Einstein's cosmological constant $\\Lambda$ with a corresponding constant energy density $\\varepsilon_\\Lambda=\\Lambda c^4/(8\\pi G)$ and negative pressure $p_\\Lambda=-\\varepsilon_\\Lambda$, assuming a positive cosmological constant. The associated cosmological models are known as $\\Lambda$CDM models.  An alternative explanation for the missing energy is {\\it quintessence}, where the dark energy density is identified with the energy density $\\varepsilon_\\phi$ (associated with a negative pressure $p_\\phi$) arising from a scalar (quintessence) field $\\phi$ (see \\cite{Peebles_Ratra_2002} for a recent review). Quintessence can be considered as a natural generalization of the cosmological constant to the case of a time-dependent $\\Lambda(t)$ with an associated time-dependent pressure.  Many cosmologists seem to accept as established that the universe is flat corresponding to $k=0$ and $\\Omega_{\\hbox{\\scriptsize tot}} := \\varepsilon_{\\hbox{\\scriptsize tot}} / \\varepsilon_{\\hbox{\\scriptsize crit}}=1$. Here we address the question: do the recent observations really establish that our universe is spatially flat? It is demonstrated that present data are {\\it consistent} with certain quintessence models possessing a constant equation of state in a hyperbolic universe, i.\\,e.\\ with negative spatial curvature, $k=-1$, corresponding to $\\Omega_{\\hbox{\\scriptsize tot}} < 1$. Our result is the consequence of the important observation that there exists a {\\it degeneracy} in the space of the relevant cosmological parameters $(w_\\phi, \\Omega_\\phi, \\Omega_{\\hbox{\\scriptsize c}})$ which are introduced below.   Our background model is the standard cosmological model based on a Friedmann--Lema\\^{\\i}tre universe with the Robertson-Walker metric \\begin{equation} \\label{Eq:metric} ds^2 \\; = \\; a^2(\\eta) \\, \\left( d\\eta^2 - \\gamma_{ij} dx^i dx^j\\right) \\hspace{10pt} , \\end{equation} where $\\gamma_{ij}$ denotes the spatial hyperbolic metric, and $a(\\eta)$ is the cosmic scale factor as a function of conformal time $\\eta$. Then the Friedmann equation reads $(a' := da/d\\eta, c=1)$ \\begin{equation} \\label{Eq:Friedmann} H^2 := \\left(\\frac{a'}{a^2}\\right)^2 \\; = \\; \\frac{8\\pi G}3 \\varepsilon_{\\hbox{\\scriptsize tot}} + \\frac 1{a^2} \\hspace{10pt} , \\end{equation} where $H=H(\\eta)$ is the Hubble parameter, and the last term in Eq.~(\\ref{Eq:Friedmann}) is the curvature term for $k=-1$. Furthermore, $\\varepsilon_{\\hbox{\\scriptsize tot}} = \\varepsilon_{\\hbox{\\scriptsize r}} + \\varepsilon_{\\hbox{\\scriptsize m}} + \\varepsilon_\\phi$, where $\\varepsilon_{\\hbox{\\scriptsize r}}$ denotes the energy density of ``radiation'', i.\\,e.\\ of the relativistic components according to photons and three massless neutrinos; $\\varepsilon_{\\hbox{\\scriptsize m}} = \\varepsilon_{\\hbox{\\scriptsize b}} + \\varepsilon_{\\hbox{\\scriptsize cdm}}$ is the energy density of non-relativistic ``matter'' consisting of baryonic matter, $\\varepsilon_{\\hbox{\\scriptsize b}}$, and cold dark matter, $\\varepsilon_{\\hbox{\\scriptsize cdm}}$, and $\\varepsilon_\\phi$ is the energy density of the dark energy due to the quintessence field $\\phi$. (In Sect.\\,\\ref{Potentials_with_Cosmological_Constant}, we will also include the energy density due to a cosmological constant.) For our later discussions of the time-dependence $\\varepsilon_{\\hbox{\\scriptsize x}} = \\varepsilon_{\\hbox{\\scriptsize x}}(\\eta)$ of the various energy components $({\\hbox{x}}={\\hbox{r}}, {\\hbox{m}}, \\phi)$, it is important to notice that the initial conditions to be imposed on the Friedmann equation (\\ref{Eq:Friedmann}) are in the case of negative curvature uniquely given by $a(0)=0$ and $a'(0)=\\Omega_{\\hbox{\\scriptsize r}}^{1/2} (1-\\Omega_{\\hbox{\\scriptsize tot}})^{-1} H_0^{-1}$ with $H_0 = H(\\eta_0)$ being the Hubble constant. Here and in the following, we use the dimensionless density parameters $\\Omega_{\\hbox{\\scriptsize x}}(\\eta) := \\varepsilon_{\\hbox{\\scriptsize x}}(\\eta) / \\varepsilon_{\\hbox{\\scriptsize crit}}(\\eta)$ with $\\varepsilon_{\\hbox{\\scriptsize crit}}(\\eta) = 3H^2/(8\\pi G)$ and $\\Omega_{\\hbox{\\scriptsize x}} := \\Omega_{\\hbox{\\scriptsize x}}(\\eta_0)$ denoting their present values.  In quintessence models, the energy density $\\varepsilon_\\phi(\\eta)$ and the pressure $p_\\phi(\\eta)$ of the dark energy are determined by the quintessence potential $V(\\phi)$ \\begin{equation} \\label{Eq:eos_phi} \\varepsilon_\\phi \\; = \\; \\frac{1}{2 a^2} \\, {\\phi'}^2 + V(\\phi) \\hspace{10pt} , \\hspace{10pt} p_\\phi \\; = \\; \\frac{1}{2 a^2} \\, {\\phi'}^2 - V(\\phi) \\hspace{10pt} , \\end{equation} or equivalently by the equation of state \\begin{equation} \\label{Eq:eos_w} w_\\phi(\\eta) \\; = \\; \\frac{p_\\phi(\\eta)}{\\varepsilon_\\phi(\\eta)} \\hspace{10pt} . \\end{equation} The equation of motion of the real, scalar field $\\phi(\\eta)$ is \\begin{equation} \\label{Eq:equation_of_motion_phi} \\phi'' \\, + \\, 2 \\frac{a'}{a} \\phi' \\, + \\, a^2 \\frac{\\partial V(\\phi)}{\\partial \\phi} \\; = \\; 0 \\hspace{10pt} , \\end{equation} where it is assumed that $\\phi$ couples to matter only through gravitation. The various energy densities are constrained by the continuity equation \\begin{equation} \\label{Eq:continuity_equation} \\varepsilon_{\\hbox{\\scriptsize x}}'(\\eta) \\, + \\, 3 (1+w_{\\hbox{\\scriptsize x}}(\\eta)) \\, \\frac{a'}a \\, \\varepsilon_{\\hbox{\\scriptsize x}}(\\eta) \\; = \\; 0 \\hspace{10pt} , \\end{equation} with the constant equation of state $w_{\\hbox{\\scriptsize r}}=\\frac 13$, $w_{\\hbox{\\scriptsize m}}=0$ for $\\hbox{x}=\\hbox{r},\\hbox{m}$ and $w_\\phi(\\eta)$ for $\\hbox{x}=\\phi$. It is worthwhile to remark that the quintessence field $\\phi$ may be regarded as a real physical field, or simply as a device for modeling more general cosmic fluids with negative pressure.  Obviously, there are two complementary approaches: \\begin{itemize} \\item[a)] Given $V(\\phi)$ compute $w_\\phi(\\eta)$, \\item[b)] Given $w_\\phi(\\eta)$ compute $V(\\phi)$, \\end{itemize} and then make predictions for or compare with cosmological observations. Among the various potentials studied in the literature (see, e.\\,g.\\ \\cite{Peebles_Ratra_2002,Brax_Martin_Riazuelo_2000} and references therein), we mention only the {\\it inverse power-law potential} \\begin{equation} \\label{Eq:potential_power} V(\\phi) \\; = \\; \\frac{A}{(B\\phi)^\\gamma} \\hspace{10pt} , \\hspace{10pt} \\gamma > 0 \\hspace{10pt} , \\end{equation} and the {\\it exponential potential} \\begin{equation} \\label{Eq:potential_exp} V(\\phi) \\; = \\; A \\, \\exp(-B \\phi) \\hspace{10pt} , \\hspace{10pt} B > 0 \\hspace{10pt} . \\end{equation} The potentials (\\ref{Eq:potential_power}) and (\\ref{Eq:potential_exp}) can be {\\it derived} \\cite{Ratra_Peebles_1988} (see also Sect.\\,\\ref{General_Properties}), if one requires a {\\it constant} $w_\\phi$ during a given evolution stage of the universe: in the radiation-dominated epoch one obtains the inverse power-law potential (\\ref{Eq:potential_power}), whereas in the quintessence-dominated epoch one requires $w_\\phi = \\hbox{const.} \\leq -1/3$ and then obtains the exponential potential (\\ref{Eq:potential_exp}).  In the following, we concentrate our attention on approach b), where one specifies the equation of state $w_\\phi$ rather than the potential $V(\\phi)$. In order to generalize the standard cosmological models like $\\Lambda$CDM models with as few as possible additional degrees of freedom, $w_\\phi$ will be chosen as constant (as studied for flat universes in, e.\\,g., \\cite{Caldwell_Dave_Steinhardt_1998,Wang_Steinhardt_1998}) This contrasts to approach a), where a whole function, i.\\,e.\\ the potential $V(\\phi)$, increases the freedom of the model enormously. Furthermore, since in the standard cosmological theories $w_{\\hbox{\\scriptsize x}}$ is constant for the various energy components, one can also consider a {\\it constant} $w_\\phi$ for the dark energy component as the most ``natural'' generalization. It turns out that the values $w_\\phi = -1/3, -2/3$ play a special role and constitute a kind of new ``exceptional phases'' in addition to the standard phases characterized by $w_{\\hbox{\\scriptsize x}} = \\{1/3, 0, -1\\}$  for $\\hbox{x} = \\{\\hbox{r}, \\hbox{m},\\Lambda\\}$. Such equations of state occur in models based on topological defects where $w = -1/3$ corresponds to a network of frustrated cosmic strings \\cite{Vilenkin_1984,Spergel_Pen_1997} and $w = -2/3$ to domain walls \\cite{Durrer_Kunz_Melchiorri_2002,Gangui_2001}.  One purpose of this paper is to derive the potential $V(\\phi)$ for the two special equations of state $w_\\phi = -1/3, -2/3$. The restriction to a constant equation of state arises from the absence of well motivated dark energy models being based on fundamental physics. This restriction should not only be considered as a prejudice but also as an approximation to time-variable equations of state when one attempts to describe cosmological observations. Models with a non-constant equation of state lead to nearly the same CMB anisotropy as corresponding models with an effective (constant) equation of state where $w_\\phi(\\eta)$ is appropriately $\\Omega_\\phi$-averaged \\cite{Doran_Lilley_Schwindt_Wetterich_2001,Doran_Lilley_2002}. As argued in \\cite{Doran_Lilley_Schwindt_Wetterich_2001}, the location of the acoustic peaks is mainly determined by $\\Omega_\\phi(\\eta_0)$, the $\\Omega_\\phi$-weighted average of $w_\\phi(\\eta)$ averaged until the present time, as well as by the average of $\\Omega_\\phi(\\eta)$ until recombination. The latter is negligible for negative $w_\\phi(\\eta)$, i.\\,e.\\ for non-tracker field models, since around $z_{\\hbox{\\scriptsize sls}}\\simeq 1100$ the energy density of the quintessence component is then subdominant compared to that of the background component. Thus $\\Omega_\\phi(\\eta_0)$ and the $\\Omega_\\phi$-weighted average of $w_\\phi(\\eta)$ suffice to approximately describe the peak structure. Since the dark energy dominates only recently, the $\\Omega_\\phi$-weighted average leads to an averaging over the recent history. Therefore, the anisotropy of the CMB is not well suited to probe the time-dependence of the equation of state $w_\\phi(\\eta)$. The main properties of the anisotropy are then determined by the angular-diameter distance $d_A$ to the surface of last scattering (see, e.\\,g., \\cite{Cornish_2000}). A further dependence of the CMB anisotropy on the dark energy arises through the late-time integrated Sachs-Wolfe effect which contributes mostly to low multipoles in the angular power spectrum. Because of the cosmic variance and a possible contribution of gravity waves, it is difficult to extract information on a time-varying equation of state \\cite{Corasaniti_Bassett_Ungarelli_Copeland_2002}. On the other hand, in classical cosmological tests based on the luminosity distance, a dark energy component contributes appreciably only for redshifts $0.2 \\lesssim z \\lesssim 2$ \\cite{Huterer_Turner_2001} since for higher redshifts the contribution to the total energy density is too minute (see also Eq.\\,(\\ref{Eq:z_m_phi})). Here the difficulties are caused by the luminosity distance which depends on $w_\\phi(\\eta)$ through a multiple-integral which smears out the information on the time-dependence of $w_\\phi(\\eta)$ \\cite{Maor_Brustein_Steinhardt_2001}. Thus, with the exception of tracker fields, it is very difficult to observe a time-dependence of $w_\\phi(\\eta)$ and therefore one can restrict the discussion to a constant equation of state.  Quintessence models with a constant equation of state differ from a cosmological constant with $w_\\Lambda=-1$ only in the value of the constant $w_\\phi\\neq -1$. If the observations are too close to $w_\\phi\\simeq w_\\Lambda$, then quintessence models might be superseded by the long-known cosmological constant, i.\\,e.\\ by the vacuum energy. Assuming {\\it flat} universes, the current bound is, e.\\,g., $w_\\phi < -0.85$ at 68\\% C.L.\\ \\cite{Bean_Melchiorri_2002} or even $w_\\phi < -0.93$ at $2\\sigma$ \\cite{Corasaniti_Copeland_2002}. The caveat is, as we shall show in Sect.\\,\\ref{Sect_Observations}, the assumed flatness. Such a value so close to the cosmological constant has, however, problems to account for the observed number of giant arcs in galaxy clusters. For a flat $\\Lambda$CDM model with $\\Omega_{\\hbox{\\scriptsize m}}=0.3$, the number of arcs is one order of magnitude too low. To obtain the right order, one either needs a low-density open model with $\\Omega_{\\hbox{\\scriptsize m}}=0.3$ or a flat model where the vacuum energy is replaced by a dark energy component with $w_\\phi \\simeq -0.6$ \\cite{Bartelmann_Meneghetti_Perrotta_Baccigalupi_Moscardini_2002}. A further variant, not discussed in \\cite{Bartelmann_Meneghetti_Perrotta_Baccigalupi_Moscardini_2002}, would be an open model with dark energy which should also produce enough strong lensing to obtain a large number of arcs. In addition, too few lensed pairs with wide angular separation are observed in comparison with the prediction of a flat $\\Lambda$CDM-model \\cite{Sarbu_Rusin_Ma_2001}. These strong lensing observations point towards a universe with negative curvature. ",
        "conclusions": ""
    },
    "0212/astro-ph0212192_arXiv.txt": {
        "abstract": "The ``Cardassian Expansion Scenario'' was recently proposed by Freese \\&  Lewis as an alternative to a cosmological constant in explaining the current accelerating universe. In this paper we investigate observational constraints on this scenario from recent measurements of the angular size of high-$z$ compact radio sources compiled by Gurvits and coworkers. We show that the allowed intervals for $n$ and $z_{eq}$, the two parameters of the Cardassian model, are heavily dependent on the value of the mean projected linear size $l$. However, the best fit to the current angular size data prefers the conventional flat $\\Lambda$CDM model to this Cardassian expansion proposal, though the latter is cosmologically credible and compatible with the $\\Theta - z$ diagram for some values of $l$. ",
        "introduction": "The standard Big Bang cosmological model is based on four cornerstones: the Hubble expansion, the Cosmic Microwave Background Radiation(CMBR), primordial Big Bang Nucleosynthesis and structure formation. Recent observations of the Hubble relation of distant Type Ia supernovae have provided strong evidence for the acceleration of the universe (Perlmutter et al. 1998, 1999; Riess et al. 1998). While current measurements of the cosmic microwave background anisotropies favor a spatially flat universe with cold dark matter \\cite{ber00,lan01}, both the deuterium abundance measured in four high redshift hydrogen clouds seen in absorption against distant quasars (Burles \\&  Tytler 1998a,b) (combined with baryon fraction in galaxy clusters from X-ray data, see White et al. 1993) and the large-scale structure in the distribution of galaxies \\cite{bah00,pea01} have made a strong case for a low density universe (for a recent summary, see Turner 2002). All these observations can be concordantly explained by the hypothesis that there exists, in addition to cold dark matter, a dark energy component with negative pressure in our universe \\cite{tur98}. The existence of this component has also been independently confirmed by other observations such as age estimates of old high-redshift galaxies \\cite{dun96,kra97,alc99} and gravitational lensing (Kochaneck 1996; Chiba \\& Yoshii 1999; Futamase \\& Hamana 1999; Jain et al. 2001; Dev et al. 2001; Ohyama et al. 2002).  During past several years, a huge number of candidates for the dark energy component have been proposed, such as a cosmological constant \\cite{wei89,car92,kra95,ost95}, a frustrated network of topological defects (such as cosmic strings or domain walls) (Vilenkin 1984; Davis 1987; Kamionkowski \\&  Toumbas 1996) and an evolving scalar field (referred to by some as quintessence) (Ratra \\& Peebles 1988; Frieman et al. 1995; Coble et al. 1997; Caldwell et al. 1998). Despite great effort to pin down the amount and the nature of the dark energy, a convincing mechanism with a solid basis in particle physics that  explains the accelerating universe is still far off. Very recently, Freese \\&  Lewis (2002) proposed a ``Cardassian Expansion Scenario'' in which the universe is flat, matter dominated and accelerating, but contains no vacuum contribution. The main point of this scenario is to modify the standard Friedman-Robertson-Walker equation as following, \\begin{equation} \\label{eq:ansatz} H^2 = A\\rho + B\\rho^n \\end{equation} where $H \\equiv \\dot R / R$ is the Hubble parameter (as a function of time), $R$ is the scale factor of the universe, and the energy density $\\rho$ contains only ordinary matter and radiation \\cite{fre02}. The second term, which may arise as a consequence of brane world cosmologies, drives the acceleration of the universe at a late epoch when it is dominent. The authors claimed that this Cardassian model survives observational tests such as the cosmic microwave background radiation, the age of the universe, the cluster baryon fraction and structure formation, and they are now studying possible observational tests \\cite{fre02}. In this paper, we give the first observational constraint on this scenario from recent measurements of the angular size of high-$z$ compact radio sources made by Gurvits et al. (1999). We show that, although this Cardassian expansion proposal is cosmologically credible and compatible with the $\\Theta - z$ diagram for some values of the mean projected linear size $l$, it is disfavoured by the best fit to the current angular size data when compared with the conventional flat $\\Lambda$CDM model. Our result is very similar to the one of Avelino \\&  Martins (2002) who used type Ia supernovae data to show another particular solution of a brane world scenario being disfavoured. There is a common point among these analysis: both models predict a universe with unreasonably low matter density. ",
        "conclusions": "Although the evidence for an accelerating universe is increasing from various astronomical observations, understanding the mechanism based on particle physics is still one of the most important challenges in modern cosmology. The ``Cardassian Expansion Scenario'' proposed by Freese \\& Lewis (2002) is one intriguing possibility, because it assumes the universe is flat, matter dominated and accelerating, but contains no vacuum contribution. We have used the updated angular size data to give the first observational constraint for this scenario. As it shown, although this Cardassian model is credible and compatible with the $\\Theta - z$ diagram for some specific values of $l$, it is disfavored by the present angular size data because of its prediction of a universe with unreasonably low matter density.  Ignorance of the characterisitic length $l$ is one of the major uncertainties in the present analysis. One method to overcome it is to do the analysis over a large enough range of $l$ to include almost all of the possibilities, and then to calculate the probability distribution for the model parameters by integrating over $l$ \\cite{che02}. However, this will loosen the cosmological constraints, making a larger range of model parameters plausible. In this sense, the Cardassian proposal will be more compatible with the present angular size data if we take into account the uncertainty of $l$.  Another uncertainty comes from the possibility that the source linear size is dependent on the source luminosity and redshift, i.e., the sources are not `true' standard rods \\cite{gur99,vis01}. Parametrizing the effects of these dependence as $l \\rightarrow l L^\\beta (1+z)^\\gamma$, Gurvits et al. (1999) and Vishwakarma (2001) have shown that the analysis with and without $\\beta$ and $\\gamma$ are basically consistent. It is reasonable, for the data compiled by Gurvits et al. has been minimized for this dependence by discarding low values of luminosities and extreme values of spectral indices \\cite{gur99,vis01}."
    },
    "0212/hep-ph0212199_arXiv.txt": {
        "abstract": " ",
        "introduction": "One of the most spectacular consequences of brane-world theories~\\cite{ADD,early,RS} with low fundamental gravitational scale, $\\mpl\\gsim\\tev$ (for some works on the structure of these theories, on their cosmological implications and experimental signatures, see \\cite{activity}), is the modification of the properties of small black holes~\\cite{admr} that, in principle, allows them to be copiously produced at the LHC and in other high-energy processes~\\cite{bf,gt,dl,support,voloshin,bhphenom,cosmic}. In a collider setting, the dominant signal for such black hole production arises from the Hawking evaporation products of the black hole, in particular the characteristic spectrum of such radiation. In a previous paper \\cite{KMR1}, we considered the leading semiclassical corrections to the spectrum of Hawking radiation for a higher-dimensional black hole evaporating to scalar modes.  We considered both a scalar field localised to the brane as well as the situation in which it could propagate in the full $(4+n)$-dimensional bulk, and computed the corresponding greybody factors as functions of energy $\\om$, angular momentum $j$, and number of (flat) extra dimensions $n$.  The scalar case is both a useful testing place of the methods used in such a computation, and potentially relevant for Higgs, or other scalar boson production from black holes formed in collisions at the LHC or other high energy colliders. However, from a practical point of view the most important cases to consider are those of the production of SM fermions or SM gauge bosons to which we turn in this paper.  In both cases the fields that need to be considered are brane localised.  For fermions this is because of the chiral nature of the SM fermions, since a fermion propagating in the bulk is necessarily non-chiral from the 4d perspective.\\footnote{The RHD neutrino can propagate in the bulk, leading to an alternate, non-seesaw explanation of the lightness of the observed neutrino states~\\cite{AHDDMR}. Thus in principle it is possible to consider the Hawking decay of a black hole to these bulk neutrino states.  However, from the perspective of the brane observer this leads only to a missing energy signal which is not distinguishable from the many other sources of missing energy in brane-world theories.  Therefore we do not consider this possibility in the present work.} It is possible to consider SM gauge fields propagating in the higher-dimensional bulk, or a subspace thereof \\cite{DDG}.  However, the already existing collider constraints bound the size of such gauge extra dimensions to $R_{\\rm gauge}^{-1} \\gsim 3\\tev$, while, in the best possible case where only the graviton propagates in the extra dimensions, the limits on the new fundamental higher-dimensional Planck scale are of order $\\mpl\\gsim 1.2\\tev$~\\cite{pdg}.   On the other hand, the validity of the semiclassical description of black holes requires that their horizon radius $r_H$ be large compared to the higher dimensional Planck length $r_H \\gg 1/\\mpl$, which given the just cited bounds, as well as the limits on the maximum mass of black hole that can be produced at the LHC, implies that $r_H$ must be somewhat large compared to $R_{\\rm gauge}$, too.  In this case the gauge fields can also be considered as effectively brane localised, the `brane' in question having some non-trivial substructure. To leading order this substructure does not change the greybody factors we compute for the spin 1 case.  Note that the properties of a black hole, in particular its Schwarzschild radius, $r_H$, and production cross section, differ from the 4d result provided that $r_H$ is smaller than the size of at least some of the extra dimensions $r_H < R$.  If the brane tension does not strongly perturb the $(4+n)$-dimensional black hole solution, and the extra dimensional space is approximately flat on the scale $r_H$, then the black hole is well-described as a angularly symmetric $(4+n)$-dimensional black hole centered on the brane, but extending out into the $n$ extra dimensions, leading to the expression~\\cite{admr,mp}, \\beq r_H = {1\\over\\mpl} \\left(M\\over\\mpl\\right)^{1\\over n+1} \\left(8 \\Gamma((n+3)/2)\\over (n+2) \\pi^{(n+1)/2}\\right)^{1/(n+1)}\\,, \\label{eq:rh} \\eeq for the Schwarzschild radius.  Here, we follow these commonly made assumptions, which have the consequence of limiting our discussion to the Hawking decay of black holes in the context of Arkani-Hamed-Dimopoulos-Dvali like theories~\\cite{ADD}.  In particular, black holes in theories of the Randall-Sundrum type \\cite{RS} require a more involved treatment (for a recent paper considering the decay of RS black holes see~\\cite{EGBK}).  After such black holes are produced, they decay by the emission of Hawking radiation, and it is expected that they will decay mostly to particles on our brane~\\cite{ehm}. The spectrum of radiation, with a temperature $T_{BH}= (1 + n)/( 4 \\pi  r_H)$ (in $G = k_B = c =\\hbar=1$ units), is given by the Hawking formula~\\cite{hawking} \\beq \\frac{dE(\\om)}{dt} = \\sum_{j,b} \\sigma_{j,b}(\\om)\\,{\\om  \\over \\exp\\left(\\om/T_{BH}\\right) \\mp 1}\\,\\frac{d^{n+3}k}{(2\\pi)^{n+3}}\\,, \\label{eq:greybody} \\eeq where $j$ labels the total angular momentum quantum number, $b$ labels any other quantum numbers of the emitted particle, as well as the particle type, and in the phase-space integral $|k| = \\om$ for a massless particle.  Here $\\si(\\om)$, known as the `greybody factor', is an energy-dependent function arising from the backscattering of part of the outgoing radiation back into the black hole by the non-trivial metric in the region exterior to the black hole horizon (for early works on the emission spectra of 4d black holes, see \\cite{classics}).  In this paper we calculate the greybody factors and emission rates for spin 1/2 and spin 1 particles localised to the SM brane in the presence of a non-rotating black hole of the type described above.  It is expected that such a non-rotating black hole is a good approximation to the exact solution after the initial `balding' and `spin-down' phases~\\cite{gt}. As discussed for the scalar case in our first paper \\cite{KMR1}, the greybody factor is computed using its equality with the absorption cross section for the appropriate type of particle incident on the background metric that describes the brane black hole. The equality of greybody factors to absorption cross sections implies that the greybody factors do not invalidate the thermal nature of the black hole.  Finally, it is important to recall that the semiclassical calculation of Hawking emission is only reliable when the energy of the emitted particle is small compared to the black hole mass $\\om \\ll M$, since only in this case is it correct to neglect the back reaction of the metric during the emission process. This in turn requires that the Hawking temperature obeys the relation $T_{BH}\\ll M$, which is equivalent to demanding that the black hole mass $M\\gg \\mpl$. Inevitably this condition breaks down during the final stages of the decay process, but for those black holes of initial mass larger than $\\mpl$ most of the evaporation process is within the semi-classical regime.  In Section~2 of this paper, we present the general equation for the radial part of the wave function of fields with spin 1/2 and 1, as well as a discussion of the metric used in calculating the greybody factors. We proceed to solve this equation at the asymptotic regimes of the near-horizon and far-field zones, for arbitrary number of extra dimensions and angular momentum number, and we match them in an intermediate zone in order to construct a complete, smooth solution over the whole radial regime.  In section~3, the greybody factors for fermionic and gauge fields are computed.  Our full analytic results are contained in Eqs.~(35-38), (40), (46), and (49).  Although analytically complicated, these full results are necessary for the accurate determination of the Hawking emission spectrum from brane black holes (see section 4).  We expect they will be of use to those searching for black holes at current and future colliders, or in ultra-high-energy-cosmic-ray collisions.  However, for ease of comparison with previous studies of greybody factors we derive simplified expressions for the leading corrections to the spectrum of the Hawking radiation valid for $\\omega r_H  \\ll 1$.  These expressions are presented in Eq. (\\ref{fermionsi}) for spin 1/2 and Eq. (\\ref{gaugesi}) for spin 1, for arbitrary $j$ and $n$. Their values for $n=2,4,6$ extra dimensions, and for the first three, and thus most important, angular momentum modes (with $j =\\frac{1}{2}$, $\\frac{3}{2}$, $\\frac{5}{2}$  for spin 1/2, and $j=$1, 2, 3 for spin 1) are tabulated in Table~1 and Table~2, respectively.  In section 4, we proceed to calculate the full greybody factors and emission rates for spin 1/2 and spin 1 fields by a numerical evaluation of our full analytic results of sections 2 and 3 (valid beyond the simplified limit $\\om r_H \\ll 1$). A similar analysis for spin 0 fields, based on the results derived in \\cite{KMR1}, is also included in this section. The results are presented in figures 1,2 and 3. Our conclusions are summarized in section 5. ",
        "conclusions": "In an earlier paper \\cite{KMR1}, we studied the problem of the decay of a higher-dimensional Schwarzschild-like black hole through the emission of either higher-dimensional bulk, or brane-localised, scalar modes. In this paper, we extended the computation of the greybody factors that appear in the Hawking decay process to the spin 1/2 and spin 1 brane-localised cases of greatest phenomenological interest. Since the angular momentum of the black hole has been ignored, our calculations are relevant to the post balding and spin-down phases of the life of a black hole produced at a high energy collider or by ultra high energy cosmic ray interactions.  The semiclassical approximation methods employed allow us to calculate the greybody factors and the corresponding emission rates for a general number $n$ of flat extra dimensions of the Arkani-Hamed, Dimopoulos, and Dvali type.  (It is important to bear in mind that black holes in highly curved bulk spacetimes, for example Randall-Sundrum 5d theories, require a somewhat modified treatment from the one presented here.)  The primary analytic results of our calculation for the greybody factors, for arbitrary $j$ and $n$, are presented in Eqs.~(35-38), and (40), (46), and (49).  Simplified expressions are provided in Eqs.~(47) and (55) spin 1/2 and spin 1 respectively. In both cases there is an increase in the greybody factor as $n$ is increased at fixed $j$. As the subsequent numerical analysis revealed, this behaviour survives up to intermediate values of the energy of the emitted particle and contributes to a considerably enhanced emission rate for both fermions and gauge bosons. A similar analysis for the emission of scalar fields also led to an enhancement of the emission rate as $n$ increases despite the fact that the greybody factor, rapidly deviating from the behaviour dictated by the leading-order correction in the low-energy approximation, is suppressed. We expect that our results will be of use in the detailed study of the signatures of possible black hole production events.   \\noindent {\\bf Note added}. While writing this paper, the related work of Ref. \\cite{park} appeared which also considers the greybody factors for black holes in brane-world theories.  \\noindent {\\bf Acknowledgments.} We wish to thank Roberto Emparan and Nemanja Kaloper for discussions."
    },
    "0212/astro-ph0212186_arXiv.txt": {
        "abstract": "We discuss the solution of accretion disk when the black hole is chosen to be rotating. We study, how the fluid properties get affected for different rotation parameters of the black hole. We know that no cosmic object is static in Universe. Here the effect of the rotation of the black hole to the space-time is considered following an earlier work of the author, where the pseudo-Newtonian potential was prescribed for the Kerr geometry. We show that, with the inclusion of rotation of the black hole, the valid disk parameter region dramatically changes and disk becomes unstable. Also we discuss about the possibility of shock in accretion disk around rotating black holes. When the black hole is chosen to be rotating, the sonic locations of the accretion disk get shifted or disappear, making the disk unstable. To bring it in the stable situation, the angular momentum of the accreting matter has to be reduced/enhanced (for co/counter-rotating disk) by means of some physical process. ",
        "introduction": "For over last three decades, fluid dynamical studies of the accretion disk around black holes have been extensively studied. Shakura \\& Sunyaev (1973) initiated this discussion considering very simplistic but effective model of accretion disk. They chose the Newtonian gravitational potential. Novikov \\& Thorne (1973) and Page \\& Thorne (1974) studied a few aspects of the accretion disk in a full relativistic treatment. Paczy\\'nski \\& Wiita (1980) proposed one pseudo-potential which could approximately describe relativistic properties of the accretion disk around non-rotating black holes. Using that pseudo-potential, Abramowicz \\& Zurek (1981) studied some transonic properties of accretion flows. Later on, with time, number of studies of accretion disk have been carried out with modified disk structures (e.g. Abramowicz et al. 1988; Abramowicz \\& Kato 1989; Chakrabarti 1990, 1996a; Narayan \\& Yi 1994; Narayan et al. 1997, 1998;  Kato et al. 1998; Mukhopadhyay 2002a) for non-rotating black holes. Also the studies of accretion disk in Kerr geometry have been explored by several authors (Sponholz \\& Molteni 1994; Chakrabarti 1996b, 1996c; Abramowicz et al. 1996; Peitz \\& Appl 1997; Gammie \\& Popham 1998; Popham \\& Gammie 1998; Miwa et al. 1998; Lu \\& Yuan 1998; Manmoto 2000). Those studies have been made either in a full general relativistic or in a pseudo-Newtonian approach.  The study on stability of accretion disk is one of the important criteria in this context. The possibility of steady and stable disk formation by incoming matter towards a black hole is allowed only for certain sets of initial parameters. Abramowicz \\& Zurek (1981) studied the effect of angular momentum on the accretion and corresponding stability of the transonic nature of the infalling matter onto the black hole. After that, Chakrabarti (1989, 1990) discussed about the formation of sonic points and shock waves in accretion disks. He showed that the formation of stable sonic points and shock in the disk are possible for a particular range of physical parameters. Recently, Mukhopadhyay \\& Chakrabarti (2001) have analysed the stability of accretion disk in presence of nucleosynthesis. As the temperature of the disk is very high (at least of the order of $10^9$K), nuclear burning can take place and generate or/and absorb energy in a large scale in the disk. Because of this generation/absorption of nuclear energy, fundamental properties of the disk, like sonic locations, temperature, Mach number, shock location (if any), etc., may be affected. Thus the stability gets influenced by the nuclear burning. If the rate of generation/absorption of nuclear energy is large, the disk structure may become unstable. In their discussion, Mukhopadhyay \\& Chakrabarti (2001) have also shown that the sonic locations disappear for a particular shell of accretion disk where the nuclear energy is significant compared to the mechanical energy of the accreting matter.  In all the above mentioned works, when the stability was analysed for the accretion disk, the central black hole was chosen to be non-rotating. As it is well known that no object is static in the Universe, before making any serious conclusions about the inner properties of the accretion disk, the consideration of rotation of a black hole is essential. Though a few works have been done for Kerr black holes, no such comparative study has been explored so far with different Kerr parameters. Thus automatically the following questions arise, does the rotation of black hole play an important role in the accretion properties of a disk? Does the consideration of rotation of black hole affect the disk structure for a particular parameter regime which is known to be stable for the non-rotating case? Here, we would like to address all these questions one by one.  Recently, we have proposed a new pseudo-Newtonian potential (Mukhopadhyay 2002b, hereinafter Paper-I) to describe the accretion disk around Kerr black holes. Following Paper-I, we will describe our accretion disk in a pseudo-Newtonian manner. The potential, that we are using, can describe approximately all the essential general relativistic properties of the accretion disk. It can reproduce the locations of marginally stable and bound orbit exactly or almost exactly for different Kerr parameters (see Paper-I), which are the pure general relativistic properties. Therefore, we call this study as a {\\it pseudo-general-relativistic approach}. As our main interest is to study the inner region of the accretion disk, we will concentrate upon the sub-Keplerian flow, where the effect of rotation of black hole is significant. In Paper-I, we already raised several questions, which are important to address in physics of accretion disk. Here we will show, how useful is the pseudo-potential prescribed in Paper-I to study the global behaviour of the accretion disk. In full general relativity, the basic equations of accretion disk are very complicated and tedious to handle. In Early, there were no suitably good pseudo-potential for a rotating black hole. It seems that Paper-I has come up with the answer and one of our aim in this paper is to show its applicability. As the location of horizon, marginally bound and marginally stable orbits change for different Kerr parameters, any inner property of the disk is very much related to the rotation of black hole. We will follow Mukhopadhyay \\& Chakrabarti (2001) to analyze the stability of the disk. We will study, whether or not the sonic location(s), shock formation(s) (if any), etc., gets affected by applying rotation to the black hole. In an accretion disk, if the sonic point looses entropy or disappears by the rotational effect of black hole, we will understand that the sonic point as well as the disk is unstable because stable sonic point is necessary for the accretion onto black hole. In the next section, we will present the basic equations needed to describe the accretion disk. In \\S 3, we will discuss about the parameter space of accretion disk and how does the rotation affect it. Subsequently, in \\S 4, we will describe the fluid dynamical results. Finally, in \\S 5, we will sum up all the results and make the overall conclusions. ",
        "conclusions": "Here, we have studied the stability of accretion disk around rotating black holes in a pseudo-Newtonian approach. Following Paper-I, we have incorporated the general relativistic effects (according to Kerr geometry) to the accretion disk in such a manner that all the essential relativistic properties, like, the locations of the marginally bound and stable orbits, black hole horizon and mechanical energy of the accreting matter are reproduced in the disk for various values of Kerr parameters. Thus we can say, our method is a 'pseudo-general-relativistic'. As our interest is to check how the rotation of a black hole solely affects the disk structure, we have chosen the inviscid fluid in accretion disk. Most of the earlier studies on structure and stability analysis of the accretion disk have been done around a non-rotating black hole (Schwarzschild geometry). As no cosmic object is static, it is very important to consider the effect of rotation of the black hole, particularly for the discussion of an inner edge of accretion disk. We have found that when the rotation of a black hole is incorporated, the valid, known disk parameter region of a non-rotating black hole is dramatically affected. As a result, the location of various sonic points and shock (if any) get shifted severely and influence the disk structure. Therefore, we can conclude that in the study of an accretion disk, the rotation should be considered. Any related conclusion depends on the rotation parameter of a black hole. So, we can say that the known physical parameter regime (Chakrabarti 1989, 1990, 1996a) to form a shock, sonic locations etc. in the disk around a non-rotating black hole is not global. The mentioned disk properties and phenomena not only get modified and/or shifted in location at the disk, but also sometimes completely disappear. In our study, we have chosen the Kerr parameter from $a=0$ to the case of an extremely high rotation as $a=0.998$ and thus, we now have a clear global picture of the accretion disk.  Though, in the literature, there are a few works on the accretion disk around Kerr black holes (e.g. Chakrabarti 1996b,c; Gammie \\& Popham 1998; Popham \\& Gammie 1998), those do not concentrate on the comparative study of a rotating and non-rotating black hole and the instability of disk induced by the rotation of black hole. Chakrabarti (1996b,c) mainly concentrated on fast rotating black holes and showed solution topology in full Kerr geometry. He does not show, how the solution varies with the change of rotation of the black hole, keeping unchanged other physical parameters. Thus, it is not clear from those works, how the Kerr parameter affects the disk solution solely. On the other hand, Gammie \\& Popham (1998) and  Popham \\& Gammie (1998) do not discuss the effect of rotation on the sonic points globally, which is shown here in Figs. 1-4. However, the fluid dynamical results achieved by them tally with the present work. For example, they (Gammie \\& Popham 1998;  Popham \\& Gammie 1998) showed, the effect of non-zero Kerr parameter is dominant close to the black hole, that is also reflected from Fig. 4 of the present work, which shows the outer sonic point branches are unchanged, while the inner ones are shifted. They also showed, the (inner) sonic points move to smaller (larger) radii as co-rotation (counter-rotation) increases in magnitude. In our work, similar features come out from Figs. 1,2,3 and 6. Thus, we can conclude once again that, the pseudo-potential prescribed in Paper-I can reproduce the general relativistic properties in accretion disk. Unlike the works by Gammie \\& Popham (1998) and  Popham \\& Gammie (1998), here we have chosen inviscid flow to understand the sole effect of rotation of the black hole on the fluid and corresponding disk structure. Finally we can conclude that the effect is significant.  Here, we have carried out the stability analysis of accretion disk by choosing the black hole to be rotating.  Prasanna \\& Mukhopadhyay (2002) have worked on the stability of accretion disk around rotating compact objects (but not for black holes) in a perturbative approach. They incorporated the rotational effect of central object in a disk indirectly, by the inclusion of Coriolis acceleration term. Here, the compact object is black hole and its rotational effect is brought from the Kerr metric only. As we have a very self-consistent pseudo-potential (Paper-I), it has been now easy to study the accretion phenomena around a rotating black hole.  From the global analysis of sonic points, we have established that the disk becomes unstable for a higher co-rotation at a particular angular momentum of the accreting matter. On the other hand, to form a shock, accretion disk must have a stable inner sonic point. Thus for higher $a$ (Kerr parameter), shock is unstable, as the inner region of disk itself is not stable. In that regime, shock may disappear by any disturbance created on the accreting matter and thus the matter may not get a steady inner sonic point to form a stable inner disk. For the counter-rotating cases, the angular momentum of the system reduces and the matter falls to the black hole more rapidly. As there is no significant centrifugal barrier to slow down the matter, the possibility of shock reduces again. For a very high counter-rotation, shock disappears completely. Thus we can conclude that the parameter region where the shock is expected to form (Chakrabarti 1989, 1990, 1996a) for non-rotating black holes, gets affected for the rotating ones. For a positive $a$, shock might have been formed for a lesser value of the angular momentum of accreting matter. On the other hand, for a negative $a$, to form a shock, the angular momentum of accreting matter should have a larger value than that of non-rotating cases. As we are studying the sub-Keplerian accretion flow, the angular momentum of matter cannot be unlimitedly high (e.g., for a non-rotating black hole, $\\lambda\\lsim 3.6742$ at last stable orbit). Thus, to form a shock, matter cannot attain the angular momentum beyond a certain limit. Also for a highly co-rotating black hole, to stabilize the disk, $\\lambda$ has to be reduced by some physical process, otherwise the radial speed could not overcome the centrifugal pressure to form a steady inner edge of the disk. Therefore, the disk would not be transonic and could be unstable in this situation. Because of a reduced $\\lambda$, centrifugal force as well as shock strength reduces, as a result, the possibility of a shock is also reduced for both the co-rotating and counter-rotating black holes. However, at an intermediate co-rotation, the parameter region of accretion disk enlarges having three sonic points and there is a possibility of a stable shock formation. On the other hand, for a counter-rotation, the valid parameter region with three sonic points at the disk shortens and thus the possibility of shock reduces.  Now the question may arise, how would a flow behave in the region where sonic points disappear because of the higher rotation of black hole? It can be thought as, matter would try to pass through same sonic point as of non-rotating (or slowly rotating) case. Because of the absence of sonic point, the searching procedure for the true sonic location should give rise to an unstable solution. Only the time-dependent simulation of accretion disk can tell whether this idea is correct or not. Overall we can say that the rotation of a black hole affects the sonic locations of an accretion disk which are directly related to the formation of stable disk structure. Therefore, the stability of an accretion disk is strongly related to the rotation of central black hole.  It has also been confirmed that the potential given in Paper-I is very well applicable for the global solution and corresponding fluid dynamical study of the accretion disk around a rotating black hole. The set of a full general relativistic disk equations is very complicated for a Kerr black hole, particularly, when the accreting fluid is chosen to be viscous. Thus, in the study of a relativistic accretion disk, this pseudo-potential is very effective to avoid the complexities of full general relativistic fluid equations. The next step should be the study of a viscous accretion disk using this general pseudo-potential."
    },
    "0212/astro-ph0212379_arXiv.txt": {
        "abstract": "We investigate the possibility of an intrinsic Baldwin Effect (i.e., nonlinear emission-line response to continuum variations) in the broad \\Hbeta\\ emission line of the active galaxy NGC 5548 using cross-correlation techniques to remove light travel-time effects from the data. We find a nonlinear relationship between the \\Hbeta\\ emission line and continuum fluxes that is in good agreement with theoretical predictions. We suggest that similar analysis of multiple lines might provide a useful diagnostic of physical conditions in the broad-line region. ",
        "introduction": "Correlations between the continuum and emission-line properties of active galactic nuclei (AGNs) afford potentially important probes of the structure and physical conditions of the broad-line region (BLR) and of the unobservable ionizing continuum in these sources. Probably the best-established correlation between continuum and emission-line properties is the ``Baldwin Effect''  (Baldwin 1977), an inverse, non-linear correlation between the equivalent width of the broad UV emission lines, most notably \\civ\\,$\\lambda1549$, and the luminosity of the underlying continuum (e.g., Osmer, Porter, \\& Green 1994). The Baldwin relationship can be parameterized as a simple power law of the form \\begin{equation} \\label{eq:basic} F_{\\rm line} \\propto F_{\\rm cont}^{\\alpha}, \\end{equation} with index $\\alpha \\ltsim 1$.  Kinney, Rivolo, \\& Koratkar (1990) and Pogge \\& Peterson (1992) investigated the effect of variability on this relationship in \\civ\\ and \\Lya. The results can be summarized as follows: \\begin{enumerate} \\item There is a {\\em global} or object-to-object relationship with slope $\\alpha \\approx 0.83$ for \\civ\\ and $\\alpha \\approx 0.88$ for \\Lya. \\item There is a rather flatter {\\em intrinsic} relationship with $\\alpha \\approx 0.4$ that describes how the line and continuum relationship changes due to variability. At least some  of the scatter in the global relationship is attributable to variability and the flatter slope of the intrinsic relationship. Moreover, the scatter in this intrinsic relationship is largely attributable to light-travel time effects within the BLR (i.e., changes in the line flux follow those in the continuum flux with some measureable time delay). \\end{enumerate}  While the Baldwin Effect is well-established in most of the strong UV lines, any global Baldwin Effect in the Balmer lines seems to be quite weak (e.g., Dietrich et al.\\ 2002). However, to the best of our knowledge, the question of whether or not there is an intrinsic relationship in the Balmer lines has not been specifically addressed. Our goal in this paper is to remedy this. We adopt an approach similar to that of Pogge \\& Peterson (1992) in attempting to remove light travel-time effects when determining the slope of the relationship. We will make use of numerous well-sampled optical spectra of NGC 5548 obtained by the International AGN Watch in this analysis to show that an intrinsic Baldwin Effect is present in \\Hbeta. ",
        "conclusions": "We have investigated the relationship between the variable \\Hbeta\\ emission line and optical continuum of NGC 5548. We discovered that an intrinsic Baldwin Effect does exist; after for correcting for light travel-time effects, we find that the amplitude of broad \\Hbeta\\ variation is less than the amplitude of variation of the underlying optical continuum. By combining these results with the relationship between the amplitudes of variability of the optical and UV continuum, we determine the relationship between the \\Hbeta\\ emission-line flux and the UV continuum flux, and find that our result is in good agreement with the predictions of photoionization calculations, even accounting for smearing of the line response by light travel-time effects within the \\Hbeta-emitting region. We suggest that similar analysis of multiple lines have potential as diagnostics of the physical nature of the BLR.  We are grateful for support of this work through NASA grant NAG5-8397. We thank Sergey Sergeev of Crimean Astrophysical Observatory for data on the recent faint state of NGC 5548. We are grateful to an anonymous referee whose comments led to an improved presentation.   \\clearpage"
    },
    "0212/astro-ph0212523_arXiv.txt": {
        "abstract": "We have developed a general purpose dust radiative transfer code for an axisymmetric system, {\\twodust}, motivated by the recent increasing availability of high-resolution images of circumstellar dust shells at various wavelengths. This code solves the equation of radiative transfer following the principle of long characteristic in a 2-D polar grid while considering a 3-D radiation field at each grid point. A solution is sought through an iterative scheme in which self-consistency of the solution is achieved by requiring a global luminosity constancy throughout the shell. The dust opacities are calculated through Mie theory from the given size distribution and optical properties of the dust grains. The main focus of the code is to obtain insights on (1) the {\\sl global} energetics of dust grains in the shell (2) the 2-D projected morphologies that are strongly dependent on the mixed effects of the axisymmetric dust distribution and inclination angle of the shell. Here, test models are presented with discussion of the results. The code can be supplied with a user-defined density distribution function, and thus, is applicable to a variety of dusty astronomical objects possessing the axisymmetric geometry. ",
        "introduction": "Astronomical systems are often surrounded by a shroud of dust. Evolved stars are the most typical of such, since they are responsible for more than 80\\% of the material annually injected into the interstellar space through dusty mass loss \\citep{sedlmayr94}. The ejected matter forms a dust-rich shell around these stars, which can be very bright in the mid-infrared (mid-IR; $\\sim 10 - 20 \\mu$m) due to thermal emission from warm (a few 100 K) dust grains. Therefore, the dust distribution in these circumstellar shells can be directly probed in the mid-IR.  Recent mid-IR observations at dust continuum have revealed toroidal density distribution in the circumstellar shells of evolved stars (e.g., \\citealt{skinner94,dayal98,meixner99}). Such axisymmetric dust distributions have been seen not only in the circumstellar shells of evolved stars but also in the shells of massive, young stars (e.g., \\citealt{ueta01b,smith02}). Most recently, the use of large aperture telescopes with mid-IR capabilities has pushed the diffraction-limited mid-IR imaging to sub-arcsecond resolution, and the intrinsically compact structure of the circumstellar dust shells has been revealed (e.g., \\citealt{jura01,ueta01a}).  However, it is not easy to interpret the mid-IR images since the mid-IR morphologies of these dust shells are highly influenced by self-extinction introduced by the geometry and inclination of the axisymmetric shells. Therefore, we need to construct numerical models in more than 1-D to properly interpret high-resolution mid-IR images of the circumstellar shells and fully understand the intertwined relationship between the morphologies and the dust distribution.  Furthermore, recent {\\iso} observations have greatly enhanced our knowledge of the circumstellar dust mineralogy (e.g., \\citealt{waters96,kemper02,molster02}). Together with the increasing availability of laboratory-measured optical constants for astronomical dust analogs (e.g., \\citealt{jager94} and subsequent series of papers; \\citealt{jager98}; \\citealt{speck98} for a recent compilation), a more elaborate treatment of dust grains is necessary to properly model the energy budget within the circumstellar dust shells, especially when dust grains are the primary means for energy transport (e.g., \\citealt{ueta01a,ueta01b,meixner02}).  In this context, we have developed a general purpose radiative transfer code, {\\twodust}, for an axisymmetric dust system. Below, we will introduce the code, mainly focusing on the treatment of dust grains (\\S \\ref{method}; also see appendices), discuss the results of test models (\\S \\ref{run}), and give a summary (\\S \\ref{2dustconcl}). ",
        "conclusions": "}  In order to numerically model a dust-enshrouded axisymmetric system, we have developed a dust radiative transfer code, {\\twodust}. This code solves the equation of radiative transfer in a 2-D polar grid by considering a fully 3-D radiation field. In this sense, the {\\twodust} code is a 2.5 dimensional radiative transfer code. A solution is sought through an iterative scheme originally developed by \\citet{collison91}, in which self-consistency of the solution is achieved by requiring a global luminosity constancy at each radial grid.  The converged solution constrains the intensity and dust temperature fields, from which we derive observables such as SED and projected surface brightness maps at given inclination angles. Dust grains in the axisymmetric system are considered to be the only means of radiative energy transport. Thus, proper considerations of the optical properties and size distributions of dust grains are of particular importance in our analysis. To calculate the absorption and scattering cross sections of dust grains, we use Mie theory supplied with laboratory-measured refractive indices of ``real'' dust grains.  We have first tested the code for a spherically symmetric shell, comparing the {\\twodust} results with the results generated by a popular 1-D radiative transfer code, {\\sf DUSTY}. The results obtained by these codes agreed quite well. Then, the scattering assumption is extended from a simple isotropic case to an anisotropic case. As expected in a spherically symmetric dust distribution, no major difference between isotropic and anisotropic cases has been recognized.  Axisymmetric models are then tested by gradually departing from spherical symmetry. We have employed a specific dust density distribution function that resulted in our investigation of post-AGB shell morphologies. Although our exploration in the parameter space is far from thorough, the results of the test models have demonstrated the relationship between the model parameters and the resulting observables (among all, the projected shell morphologies). As in the spherical cases, the assumptions of isotropic and anisotropic scattering are both tested with the axisymmetric models. The test results are consistent with the expectations that the presence of large grains would bring the anisotropically scattered light farther away from the source of scattered radiation based on the known forward/backward scattering tendency of dust grains.  These tests have demonstrated the basic capabilities of this new code fairly well. No spurious result has been generated in the entire test model runs, and therefore, we are reasonably confident to conclude that the {\\twodust} code would produce good dust radiative transfer models for an axisymmetric system. With these model runs, we have also been able to characterize the effects of various parameters involved in the model to the model results, especially the resulting SED and shell morphologies. Although such characterizations are far from complete, this is nonetheless a necessary and important first step to confidently apply the {\\twodust} code in the following model fitting using observed data. In fact, the {\\twodust} code has been applied to model the post-AGB shells and have successfully produced the best models to date to suggest that the two distinct types of post-AGB shell morphologies arise mainly from the difference in the optical depth of the shells \\citet{meixner02}.  We also demonstrated the difficulty of constraining the shell geometry and inclination angle solely from the shape of the SED, since spatially unresolved SEDs do not provide any spatial information necessary to constrain the geometry and inclination angle of the dusty shells/disks. Such spatial information needs to be obtained from high-resolution images. To investigate the energetics and geometry of the dusty shells/disks, especially important are mid-IR images that directly probes the dust distribution within the shells/disks."
    },
    "0212/astro-ph0212073_arXiv.txt": {
        "abstract": "{ Deep exposures of the metal-rich globular clusters NGC\\,6496\\ and NGC\\,6352 were obtained with the WFPC2 camera on board the Hubble Space Telescope (HST) through the F606W and F814W filters. The resulting colour-magnitude diagrams (CMD) reach down to absolute magnitude $M_{814} \\simeq 10-10.5$, approximately 5 magnitudes below the main sequence (MS) turn-off (TO). The MS of the two clusters are sharp and well defined and their fiducial lines overlap almost exactly throughout this range. Their colour is, however, more than $0.1$\\,mag redder than the MS fiducial line of the prototype metal-rich globular cluster 47\\,Tuc (NGC\\,104), after proper correction for the relative distances and reddening. This provides solid empirical evidence of a higher metal content, which is not surprising if these objects belong indeed to the bulge as their present location suggests.  A good fit to the upper part of the MS of both clusters is obtained with a $10$\\,Gyr-old theoretical isochrone from Baraffe et al.  (1998) for a metallicity of [M/H]=$-0.5$, but at lower luminosities all models depart considerably from the observations, probably because of a deficiency in the treatment of the TiO opacity.  The luminosity functions (LF) obtained from the observed CMD are rather similar to one another and show a peak at $M_{814}\\simeq 9$. The present day mass functions (PDMF) of both clusters are derived down to $M_{814}\\simeq 10.5$ or $m \\simeq 0.2$\\,M$_\\odot$ and are consistent with power-law indices $\\alpha=0.7$ for NGC\\,6496 and $\\alpha=0.6$ for NGC\\,6352. The PDMF of NGC\\,104 is twice as steep in the same mass range ($\\alpha=1.4$). We investigate the origin of this discrepancy and show that it can be understood if the two clusters contain a considerably higher fraction of primordial binaries amongst their MS population, similar to that expected in the bulge.  We briefly discuss the implications of this finding on the process of star and binary formation and on the universality of the IMF. ",
        "introduction": "Since the work of Zinn \\& West (1984), Armandroff \\& Zinn (1988) and Armandroff (1989), the disc population of the Galactic globular clusters (GC) has been recognised as a different subsystem, both dynamically and chemically, compared with that of the halo clusters. A detailed knowledge of these subsystems is important for our understanding of the fundamental properties of GC such as the concentration parameter, the shape of the mass function (MF), their age and chemical composition, since they correlate with the Galactocentric distance (Meylan \\& Heggie 1997) and/or height above the Galactic plane.  Studying the disc GC, although difficult because of the foreground star contamination and high extinction due to their low Galactic latitude, could shed light on the disruption mechanisms by dynamical friction and bulge shocks, which deeply modify their MF in time. Furthermore, investigating the physical differences between their stellar population and that of the halo clusters could improve our understanding of the formation process of our Galaxy.  In this paper, we present the analysis of the deep CMD of NGC\\,6496 and NGC\\,6352, two clusters in the direction of the bulge and putative members of this subgroup, based on observations collected with the WFPC2 on board the HST.  NGC\\,6496 was originally believed to be a member of the disc system of GC, but Richtler et al. (1994) questioned this classification. They suggested instead that NGC\\,6496, together with two other clusters, NGC\\,6624 and NGC\\,6637, could be halo clusters with strongly inclined orbits. NGC\\,6496 lies in the Southern sky at $\\alpha_{2000}=17^{o}59\\arcmin02\\arcsec$ and $\\delta_{2000}=-44^{0}15\\arcmin54\\arcsec$. Its Galactic coordinates are $l=348.02^{o}$ and $b=-10.1^{o}$. A distance of $R_G=4.3$\\,kpc from the Galactic centre and $Z_G=-2.0$\\,kpc below the Galactic plane (Harris 1996), place this object at the edge of the bulge.  Notwithstanding the fact that the surrounding field shows no strong stellar foreground variation, its low Galactic latitude has hampered the observation of this cluster from the ground. Indeed, most of the contradictory estimates as to its metallicity, reddening, and Galactocentric distance probably arise mainly from the difficulty of observing objects through the dusty and crowded Galactic disc.  Armandroff (1988) presented the first CMD of NGC\\,6496 in which the photometry reached $\\sim 2$ magnitudes below the horizontal branch (HB), disclosing for the first time the usual red arm of the metal-rich clusters. The extinction towards NGC\\,6496 is uncertain, with estimates ranging between $E(B-V)\\,=\\,0.09$ and $E(B-V) = 0.24$ (Armandroff 1988; Zinn \\& West 1984; Burstein \\& Heiles 1982; Richtler et al. 1994).  The metal content of NGC\\,6496 has also been controversial for a long time. Friel \\& Geisler (1991), in their work on the disc GC based on Washington photometry, used the reddening value of $0.19$ by Armandroff (1988) and derived for NGC\\,6496 a metallicity of [Fe/H]$=-1.05$. Richtler et al. (1994) measured its metallicity on the basis of the morphological characteristics of the CMD, obtaining [Fe/H]$=-1.0$. Their value is in agreement with the metallicity coming from Washington photometry, but substantially lower than that given by Zinn \\& West (1984), i. e. [Fe/H]$ = -0.48$, and adopted by Armandroff (1988). The discrepancies in the metallicity scale amongst different authors might be explained by the intimate connection between reddening and metallicity estimates.  NGC\\,6352 is located at $\\alpha_{2000}=17^{o}25\\arcmin29\\arcsec$ and $\\delta_{2000}=-48^{o}25\\arcmin22\\arcsec$. Its Galactic coordinates (${\\it l}=341.4^{o}$, ${\\it b}=-7.2^{o}$) and distance from the Galactic centre and plane ($R_G=3.3$\\,kpc, $Z_G=-0.7$\\,kpc; Harris 1996) place this objects at the edge or within the bulge. Based on this information and on radial velocity and metallicity measurements, NGC\\,6352 has been recognised, since Zinn (1985) and Armandroff (1989), as an unambiguous member of the disc group of globular clusters. Sarajedini \\& Norris (1994) presented CCD ground-based photometry for NGC\\,6352 reaching $\\sim 2$ magnitudes below the level of its HB, disclosing the red clumped morphology typical of metal-rich clusters. By simultaneous estimate of both metallicity and reddening, using theoretical stellar evolutionary tracks, they obtain $[Fe/H]=-0.51$, in the Zinn \\& West metallicity scale, and $E(B-V)=0.24$.  Fullton et al. (1995) observed NGC\\,6352 with the WFPC1 on board HST, obtaining a CMD reaching $\\sim 3$ magnitudes below the MS TO. Their main results are that NGC\\,6352 is slightly more metal rich than 47\\,Tuc and that it has the same age as 47\\,Tuc within the errors of their photometry. These results imply that the primordial gas might have collapsed in a disc configuration at an early epoch of the formation of our Galaxy.  In light of all the uncertainties affecting the determination of the extinction, distance and metallicity of these clusters, the skeptical reader might wonder as to how meaningful and practically useful it is to classify these objects as being members of one or another subgroup of the GC. We believe that the availability of new, powerful instruments such as the HST and VLT should call for a critical re-evaluation of the semi empirical methods on which such classifications were based.  In the following sections we provide solid observational data in the form of high precision CMD of NGC\\,6496 and NGC\\,6352, covering almost the full range of stellar masses from the TO to $\\sim 0.2 \\msun$. We compare them with the CMD of the archetype, well-studied metal-rich globular cluster 47\\,Tuc (NGC\\,104) and address the issue of their metallicity in a way so far never attempted, yet robust in light of its simplicity.  Furthermore, from the LF we derive the best fitting underlying PDMF to test the hypothesis as to whether the IMF could be metal dependent and, therefore, differ here from that of the metal-poor halo clusters studied by Paresce \\& De Marchi (2000). ",
        "conclusions": "To the expert eye, this discrepancy would probably not come as a surprise, since it is the natural consequence of the already different LF seen in Figure\\,\\ref{fig6}. To be precise, however, the difference in the LF per se would not necessarily imply that also the underlying MF should be different since 47\\,Tuc has a lower metal content than the other two objects. Thus, stars of equal mass would appear slightly brighter in 47\\,Tuc due to the different M--L relation, therefore changing the shape of the observed LF. In fact, it is well understood that the most notable feature in the LF of a cluster, i.e. the brightness at which it peaks, shifts to fainter magnitudes with increasing metallicity (see e.g. D'Antona 1998; Kroupa \\& Tout 1997; von Hippel et al. 1996; De Marchi \\& Paresce 1995b) to the point that, if the distance to a cluster were known, the magnitude of the peak could be used as a metallicity indicator.  The application of the M--L relations appropriate to each cluster has, however, removed the signature of the metallicity. As explained in Section\\,5, although some uncertainties exist as to the exact shape of the M--L relationship in the range $-0.5 < {\\rm [Fe/H]} < 0.0$ (Baraffe et al. 1998), they cannot be held responsible for the large discrepancy in the MF seen in Figure\\,\\ref{fig3}. We must, therefore, accept that NGC\\,6352 and NGC\\,6496 have MF which are inherently different from that of 47\\,Tuc and from the dozen lower metallicity halo clusters studied to date in detail with the HST (Paresce \\& De Marchi 2000).  The effects of energy equipartition and the resulting mass segregation have been observed and studied in detail in several clusters and can be modelled satisfactorily (47\\,Tuc, Paresce et al. 1994; NGC\\,7078, De Marchi \\& Paresce 1996; NGC\\,6397, King et al. 1995, De Marchi et al. 2000; NGC\\,6121, Pulone et al. 1999; NGC\\,6656, Albrow et al. 2001). It is today understood that, near the cluster's half-mass radius, the local MF faithfully reflects the global MF of the whole cluster (as originally suggested by Richer et al. 1991). Since the observations of NGC\\,6352 and NGC\\,6496 on which we here report were conducted at their half-mass radius (see Section\\,2), the effects of internal dynamics can be safely excluded from the picture.  On the other hand, the action of the Galactic tidal field can be equally strong in reshaping the MF of a cluster through disc and bulge shocks during its repeated passages at the perigalacticon. For example, the MF of NGC\\,6712 measured near the half-mass radius (De Marchi et al. 1999; Andreuzzi et al.  2001) bears no resemblance to that of any other cluster studied so far, in that if falls steeply with decreasing mass. Pal\\,5 (Grillmair \\& Smith 2001) and E\\,3 (van den Bergh et al. 1980) might have undergone a similar, although less severe fate.  It is, thus, fair to wonder whether the MF of NGC\\,6352 and NGC\\,6496 were similarly affected by the Galaxy.  Gnedin \\& Ostriker (1997) and Dinescu et al. (1998) used the mass, size and space motion parameters of most Galactic globular clusters to infer the magnitude of their past dynamical interaction with the potential field of the Galaxy. They define a useful parameter, the time to disruption ($t_d$), which should directly correlate with the strength and extent of the overall tidal interaction and, thus, ultimately with the shape of the present day MF. However, whilst it is true that for NGC\\,6712 they predict a probability of disruption about two orders of magnitude higher than average, Paresce \\& De Marchi (2000) have shown that no correlation exists between $t_d$ and the shape of the MF of the 12 clusters that they studied in detail with the HST. The value of $t_d$ must thus be used with care and we shall do so here to show, with a very simple reasoning, that neither NGC\\,6352 nor NGC\\,6496 can have been severely affected by tidal stripping. If the theory of Gnedin \\& Ostriker (1997) is indeed correct, the destruction rate for NGC\\,6496 is about 300 times larger than that of NGC\\,6352, yet the two clusters have the same global MF! So, if we want these clusters to have had an IMF originally as steep as that of 47\\,Tuc, we cannot hold tidal stripping accountable for the alleged flattening. Moreover, postulating that the two IMF were originally different from one another but that, through dynamical evolution, they have been brought into agreement would look highly contrived, in light of the widely different dynamical histories of the two clusters. Lacking any other viable explanation, it is more logical to conclude that the IMF of these clusters did not change with time, i.e.  that it was already flatter at birth and the same for both.  This conclusion, however, does still not explain why the IMF of NGC\\,6352 and NGC\\,6496 should be flatter than that of the less metal rich GC unless, of course, it is the metallicity itself that governs the shape of the IMF. Claims of an apparent correlation between the shape of the MF and cluster metallicity have been put forth in the past by McClure et al. (1986), who analysed a set of nine, relatively deep GC LF obtained from the ground and noticed a shallower slope for richer clusters, in the same sense that we report here. More recent data, however, do not confirm these results (Paresce \\& De Marchi 2000). Von Hippel et al. (1996) went even further and studied how metallicity affects the shape of the IMF of clusters of various ages (from young open cluster to globulars), concluding in favour of an invariant IMF.  Whether stars form through the hierarchical fragmentation of a proto-stellar cloud (Elmegreen 1999) or through the accretion of smaller clumps (Bonnell et al. 2001), once the pre-MS phase begins, the metallicity and ensuing atmospheric opacity will eventually determine the limiting mass for stable Hydrogen burning (see e.g. D'Antona \\& Mazzitelli 1996) and, thus, the value of the IMF peak mass. In other words, the metal content of the natal cloud will affect its temperature and pressure, on which the characteristic mass eventually depends (Larson 1998). If turbulent fragmentation prevails below $\\sim 1$\\,M$_\\odot$, its dependence on the cooling rate and, hence, on the metallicity could affect the shape of the IMF, but no clear prediction as to the results of this process exists (Larson 2002). Moreover, the observational evidence seems overwhelming that the IMF is universal and independent of the total mass, density, age and metallicity of the stellar population (Scalo 1998; Gilmore 2001; but see also the cautionary remark in Kroupa 2001).  NGC\\,6352 and NGC\\,6496 would seem, at least apparently, to violate this rule. We have, however, so far ignored the important contribution of binaries (Kroupa 2001). Other examples exist of relatively old and metal rich systems that display a rather flat MF: the bulge (Holtzman et al. 1998; Zoccali et al. 2000), the spheroid (Gould et al. 1997) as well as the Galactic disc in the solar neighbourhood (Kroupa et al. 1993; Ried \\& Gizis 1997; Gould et al. 1997). In all cases, the MF of these systems, as derived directly from the observed LF, is rather flat if no account is taken of the presence of binaries: for instance, for stars less massive than $\\sim 0.5$\\,M$_\\odot$ Gould et al. (1997) find $\\alpha = 0.56$ in the disc and $\\alpha = 0.75$ in the spheroid. Ignoring the binaries, however, results in the underestimate of the real IMF slope by an amount that depends on the slope itself and on the binary fraction (Sagar \\& Richtler 1991; Kroupa 2001) and which can be of order $\\Delta\\alpha \\simeq 0.5$ or higher. As Holtzman et al. (1998) show, a measured MF slope of $\\alpha\\simeq 0.7$ can be in reality as steep as $\\alpha\\simeq 1.3$ if the binary fraction is of order 50\\,\\%.  The age and metallicity of the stellar population in NGC\\,6496 and NGC\\,6352 should be rather similar to those of the stars in the bulge, as one can judge for instance from the similarity of their CMD and particularly of their curved red giant branch and clumped horizontal branch (Richtler et al. 1994; Fullton et al. 1995; Zoccali et al. 2000). It is, then, reasonable to apply to their MF a correction similar to that derived by Holtzman et al. (1998), which would in turn bring them into good agreement with that of 47\\,Tuc and of the other, lower metallicity halo clusters (Paresce \\& De Marchi 2000).  At this juncture, one might wonder whether a similar correction for binaries should be applied to all GC, including 47\\,Tuc. Albrow et al. (2001), however, have recently conducted high precision differential time series photometry on $\\sim 46,000$ stars in the core of this cluster and have concluded that the binary fraction is on average $\\sim 13\\,\\%$ within $\\sim 4$ core radii ($r_c$), with indications that it drops to $\\sim 8\\,\\%$ or less outside of $2.5\\,r_c$. A correction to the IMF in this case is, therefore, not necessary. The only notable signature of binaries on the observed CMD is a vertical broadening of the MS (of up to $0.75$\\,mag along the magnitude axis for equal mass systems), which should be best observed where the MS slope is not too steep in the CMD, or where it changes abruptly near the MS kinks due to changes in the main source of opacity (see, e.g., the CMD of NGC\\,6752 by Rubenstein \\& Bailyn 1997). With the advent of high precision HST photometry, this test has become possible (De Marchi \\& Paresce 1995b; Elson et al. 1995; Cool et al. 1996) and has revealed the paucity of these objects. Yet, in only a few cases, for example NGC\\,6752 (Rubenstein \\& Bailyn 1997), has the photometry been sufficiently precise to detect this binary sequence.  The MS of NGC\\,6496 and NGC\\,6352 is, however, much steeper than that of NGC\\,6752 in the magnitude range where the photometric uncertainty would allow us to see the effects of a broadening. Below $m_{606}\\simeq 23$ the MS of NGC\\,6496 flattens out, but here the photometric uncertainty dominates (see Table\\,2). Furthermore, since the vertical displacement of a binary with respect to the MS in the CMD depends on the mass ratio $q=m_2/m_1$, it will be easier to detect nearly equal mass systems, which are however believed to be a small fraction of the total, at least according to Duquennoy \\& Mayor (1991). These authors investigate the multiplicity of nearby G dwarfs and find that the secondary mass distribution peaks at around $q\\simeq 0.25$, where the solar neighbourhood's IMF has its maximum (Kroupa et al. 1993), thereby suggesting that binaries should form some time after the formation of single stars.  If the same mass distribution applies to the binaries in NGC\\,6496 and NGC\\,6352, their signature on the MS will remain undetected, at least at the present level of photometric accuracy, since a typical $\\sim 0.25$\\,M$_\\odot$ companion would add insignificantly to the brightness of a more massive primary. Although for late M dwarf primaries the $q$ distribution seems somewhat flatter than for G stars (Fischer \\& Marcy 1992), again in agreement with a binary formation scenario through random sampling from the same IMF, our photometric uncertainty, the contamination due to field objects and the steepening of the MS in the CMD plane prevent us from reliably probing this region of the parameter space.  Whilst NGC\\,6352 and NGC\\,6496 are the only two clusters in the bulge for which a deep MF has been measured so far, at this stage one could tentatively conclude that the main effect of the metallicity difference between these and the halo clusters may indeed be the binary frequency. The proposed universality of the IMF would not seem to be affected by this conclusion, but it must not be disconnected from the process of binary formation. In order for the final stellar IMF to be the same regardless of the metal content, it would require that stars are all born single and that only later they aggregate in multiple systems within the natal cloud (where they are observed), in a way which however depends on the metallicity. That the properties of binaries can be better explained through a ``two step'' IMF has been discussed recently by Durisen et al. (2001). Thus, it cannot be excluded, for instance, that the stronger UV flux and more powerful winds of OB stars in metal poor clusters would not only hinder the formation of low mass objects (Larson 1998), but also discourage their interaction and aggregation into multiple systems. This scenario agrees with the finding (Duquennoy \\& Mayor 1991; Fischer \\& Marcy 1992) that, in the solar neighbourhood, binary formation seems to happen through random association of two stars drawn from the same IMF. An alternative proposal is also possible in which, after the initial collapse, clumps form which further fragment into single stars or multiple systems depending on the cloud's metallicity. This hypothesis would, however, be less favoured because in order to conserve the uniformity of the stellar IMF with metallicity the fragmentation mechanism should adapt itself to the metal content in a way which is presently not understood.  In spite of the rather speculative nature of the previous paragraph, we stress here that our result is the first tentative indication of a large fraction of binaries in some globular clusters."
    },
    "0212/astro-ph0212559_arXiv.txt": {
        "abstract": "There is evidence that cosmic rays were present in galaxies at moderately high redshift. This suggests that magnetic fields were also present. If cosmic rays and magnetic fields must always be close to equipartition, as they are to an order of magnitude within the local universe, this would provide a powerful constraint on theories of the origin and evolution of magnetic fields in galaxies. We evaluate the role of magnetic fieldstrength in cosmic ray acceleration and confinement. We find that the properties of small scale hydromagnetic turbulence are fundamentally changed in the presence of cosmic rays. As a result, magnetic fields several orders of magnitude weaker than present galactic fields can accelerate and retain a population of relativistic cosmic rays, provided that the fields are coherent over lengthscales greater than a cosmic ray gyroradius. ",
        "introduction": "Superthermal particles, or cosmic rays, are a major constituent of the interstellar medium in galaxies. They are dynamically coupled to the thermal gas, and collectively provide pressure support, can drive outflows, and are thought to modify the structure of shocks. Cosmic rays are the dominant source of ionization, and a major source of heating, in regions of high visual extinction. Spallation reactions between cosmic rays and ambient interstellar nuclei are the primary mechanism for synthesizing light elements.  The energy spectrum, directional anisotropy, and chemical composition of cosmic rays are measured directly through \\textit{in situ} observations on the Earth and in the heliosphere. The global properties of cosmic rays in the Galaxy are probed through their radio frequency synchrotron emission and $\\gamma$-ray line and continuum emission. A recent summary of the observations is given by \\citet{S02ii}.  Although synchrotron and $\\gamma$-ray radiation from normal galaxies have up to now been directly detected only within the local universe, there is evidence that cosmic rays existed in galaxies even at early times. The light elements present in metal poor halo stars are thought to originate by spallation reactions between energetic particles and the material from which these stars formed [\\citet{DLL92}, \\citet{G92}, \\citet{D98}, \\citet{B99} and references therein]. Further evidence is provided by the spectrum of the diffuse $\\gamma$-ray background, which is best fit by a superposition of active and normal galaxy $\\gamma$-ray spectra at moderate redshift [\\citet{VF02}].  The energy densities in the Galactic magnetic field and cosmic rays are now roughly equal. If this relationship is universal and necessary, then the presence of cosmic rays in young galaxies implies the existence of equipartition magnetic fields. This would be an important constraint on theories of the history of magnetic fields in galaxies, a subject on which evidence is sparse. At present, the only direct observations of magnetic fields associated with young galaxies are the Faraday rotation measurements in damped Ly${\\alpha}$ systems at redshift $z\\le 1$ reported by \\citet{OW95}. These observations suggest magnetic fields of $\\sim 1 \\mu$G strength which are spatially coherent over several kpc, similar to the fields in the Milky Way and other spiral galaxies. Inferences based on cosmic rays could probe galactic magnetic field evolution at earlier epochs, and, we will see,  could be directly relevant to conditions in star forming regions. This is significant because of the important role magnetic fields can play in star formation and their possible effect on the Initial Mass Function.  The leading theories of cosmic ray acceleration and confinement depend on the strength of the magnetic field. In this paper we use this dependence to estimate the minimum magnetic field at which a population of relativistic cosmic rays with energy density comparable to the present Galactic population can be accelerated and confined. We find that a field several orders of magnitude weaker than the present Galactic field is sufficent for both acceleration and confinement, provided that the coherence length of the field exceeds the cosmic ray gyroradius.  The plan of this paper is as follows. In \\S\\ref{sec:early} we state our assumptions about those conditions in early galaxies which affect the acceleration and propagation of cosmic rays. In \\S\\ref{sec:review} we briefly review the theory of energetic particle transport and its application to cosmic rays. In \\S\\ref{sec:dmin} we derive constraints which arise by assuming that the cosmic rays diffuse at the minimum possible rate, which at a given fieldstrength gives the maximum acceleration and confinement. In \\S \\ref{sec:highbeta} we discuss hydromagnetic turbulence in a weak magnetic field and find that it is drastically affected by even a small population of cosmic rays. Fluctuations grow rapidly in the presence of cosmic ray anisotropy, and, in contrast to fluctuations driven by cosmic ray anisotropy in the present Galactic environment, are not Alfv\\'enic in character. At small amplitudes, the damping rate is much lower than the excitation rate. Therefore, we predict that the waves grow to nonlinear amplitudes. This may have implications for the nature of turbulence in other environments in which the magnetic fields is relatively weak and a population of energetic particles is present. Section \\ref{sec:discussion} is a summary and conclusion. Most of the technical material is in Appendices A - C.  We use Gaussian cgs units throughout, except in expressing the energies of cosmic rays, where we follow standard practice and use eV, and also in certain formulae, in which we use ``natural\" units such as parsecs or years. ",
        "conclusions": "}  The light elements present in the oldest Galactic halo stars, and the spectrum of the diffuse $\\gamma$-ray background, are indirect evidence that cosmic rays were present in young galaxies. Cosmic rays cannot be accelerated or confined without some level of magnetization. In this paper, we investigated just how strong a magnetic field is required, and what the presence of cosmic rays implies about the evolution of galactic magnetic fields.  Cosmic rays can be confined in appreciable numbers only if the magnetic coherence length $L^B$ is larger than their gyroradii. This is the standard argument for the extragalactic origin of the cosmic rays with energies above $\\sim 10^9$ GeV. Generalizing the argument to early times leads to a lower limit on the product $BL^B$ [eqn. (\\ref{equ:rg})]. If $L^B$ were as large as a kpc, typical of the gradient lengthscale in the galaxy itself, then $B$ must have exceeded 10$^{-15}$G if GeV protons were confined. If the magnetic energy density was concentrated at a much smaller scale such as the the Ohmic dissipation length, cosmic rays probably could not have been confined at all.  Under contemporary conditions, cosmic rays are scattered by gyroresonant interactions with small amplitude ($B_1/B_0\\sim 10^{-3} - 10^{-4}$) hydromagnetic turbulence [eqn. (\\ref{equ:resonance})], and propagate diffusively. Although the turbulence can in principle be excited by the streaming anisotropy of the cosmic rays themselves, it is also strongly damped, and it appears that supplemental sources of turbulence are required\\footnote{If the turbulence is highly anisotropic, so that the efficiency of scattering is reduced, then processes other than gyroresonant scattering, such as magnetic mirroring, may be necessary to explain confinement \\citep{C00a}}. The excitation of waves by cosmic rays is, however, important in the vicinity of strong shock waves, where it is believed that the bulk of galactic cosmic rays are accelerated.   If cosmic rays in early galaxies propagated diffusively, the limits on the fieldstrength are more stringent.  We derived lower bounds on the fieldstrengths necessary for diffusive shock acceleration and for confinement by assuming that the scattering frequency is the gyrofrequency. This corresponds to a nonlinear level of turbulence [$B_1/B_0\\sim 1$; eqn. (\\ref{equ:nu})]. We used the upper bound on the acceleration rate derived by \\citet{LC83} to estimate the minimum magnetic fieldstrength required to accelerate protons to relativistic energies [eqns. (\\ref{equ:Eb}), (\\ref{equ:Es}), (\\ref{equ:Esb}), and (\\ref{equ:Ed})] in supernova and superbubble driven shocks, and found it to be about 10$^{-10}$ G. This turns out to be similar to the minimum fieldstrength required to confine cosmic rays to superbubbles long enough for them to undergo nuclear collisions with ambient material and synthesize the light elements observed in the oldest stars [eqn. (\\ref{equ:bubbleconf})]. At this fieldstrength, the growth rate of the streaming instability is fast enough that the cosmic rays may be able to trap themselves in the vicinity of the shock front [see discussion following eqn. (\\ref{equ:rootsnumerical})].  The magnetic fieldstrengths in regions of active star formation may not have been representative of the global galactic fieldstrength (see the discussion in \\S\\ref{sec:early}). Therefore we derived a relationship between the average interstellar energy density in cosmic rays, the power in cosmic ray sources, and the global confinement time, assuming the maximum scattering rate. If the intensity of cosmic ray sources and the energy density in cosmic rays were comparable to their current values, the magnetic field could have been as low as 10$^{-12}$ - 10$^{-13}$ G. The fieldstrength scales inversely with the luminosity of the sources, so an enhanced supernova rate would have permitted an even smaller magnetic field. A field of this strength is consistent with growth of the waves on a timescale less than the nominal confinement time of 10$^{7}$ yr, provided that the drift velocity $v_D$ exceeds $\\sim$100 km s$^{-1}$.  Estimates of the fieldstrength based on cosmic ray diffusion theory scale with the turbulent amplitude as $(B_1/B_0)^{-2}$ [see eqn. (\\ref{equ:nu})]. There are two reasons why the turbulence could have been much stronger at early times than it is today. The sources extraneous to cosmic rays range from stellar radiative and mechanical luminosity to large scale dynamical effects associated with differential rotation and self gravity [\\citet{SB99}, \\citet{KO00}]. Under present conditions, the gyroradius of a GeV cosmic ray is about 1 AU [see eqn. (\\ref{equ:k0})], much smaller than the scale at which turbulent energy is injected. But at fieldstrengths of 10$^{-12}$ G, the gyroradii were at parsec scales, closer to the energy injection scale, at which the turbulent amplitude is significantly larger.  The strength of these turbulent sources, however, can only be speculated upon. Turbulence is driven by the anisotropy of the cosmic rays themselves. The anisotropy is an inevitable consequence of the discrete and inhomogeneous distribution of cosmic ray sources, and the finite size of the galaxy. In \\S\\ref{sec:highbeta} we considered the propagation, excitation, and damping of low frequency electromagnetic waves in a medium consisting of cosmic rays, thermal plasma, and a weak magnetic field. We found that the circularly polarized, parallel propagating Alfv\\'en waves which efficiently scatter cosmic rays in contemporary galaxies, and which are excited by cosmic ray streaming anisotropy, are dramatically modified because the electrons $\\bE \\times\\bB$ drift in the wave electromagnetic fields but the cosmic rays do not [Appendix A and eqn. (\\ref{equ:Xdef})]. These modifications are only significant when the pressure in the mean field is much less than the cosmic ray pressure [see eqn. (\\ref{equ:Xdef})]. We found that if the drift velocity $v_D$ is far above the instability threshold $\\vai$, the real and imaginary parts of the wave frequency are comparable, and much larger than the Alfv\\'en frequency [see eqn. (\\ref{equ:rootslimitnew})]. In contrast to the current situation, neither Landau damping in fully ionized regions (\\S\\ref{subsec:landau} and Appendix B) nor frictional damping in weakly ionized regions (\\S\\ref{subsec:friction} and Appendix C) appears capable of competing with the excitation rate. Therefore, we predict that the waves grow to nonlinear amplitudes, and the cosmic rays are scattered at nearly the maximum possible rate. The limits on $B$ obtained under the assumption of scattering at the maximum rate are probably close to correct.  Cosmic rays will interact dynamically with the thermal gas as long as they are strongly scattered [see eqns. (\\ref{equ:Fcr}) and (\\ref{equ:dotEcr})]. The rate of energy transfer is usually assumed to be proportional to the Alfv\\'en speed under standard conditions, but is much larger in the present case because the turbulence which scatters the cosmic rays is highly super-Alfv\\'enic [see eqn. (\\ref{equ:vphivai})]. Cosmic rays could potentially play an important role in driving outflows from young galaxies.  The fact remains that magnetic fields and cosmic rays are in approximate equipartition in our own Galaxy at the present time. We suggest that this is so because they share a common energy source: supernovae. Given enough time, magnetic fields, cosmic rays, and for that matter turbulent bulk motions will probably achieve a state near equipartition, as currently observed in the Galaxy. However, as we have shown in this paper, this is not required by anything fundamental to the physics of cosmic ray acceleration or confinement. Why the magnetic field and cosmic ray energy density have saturated at values close to the turbulent energy density, instead of increasing beyond it, is not completely clear. The reason may be that if the magnetic field or cosmic ray energy density were appreciably larger, the vertical stratification of the interstellar medium would be unstable to buoyancy driven instabilities \\citep{P66}, which would probably lead to separation of the thermal and nonthermal components of the interstellar medium, and possibly facilitiate magnetic field and cosmic ray escape.  We have briefly considered the role of these buoyancy instabilities at early times, and find it to be ambiguous. If the ratio of cosmic ray pressure to thermal gas pressure is denoted by $\\alpha$, the effective polytropic exponents of the thermal gas and cosmic rays are $\\gamma_g$ and $\\gamma_{cr}$, and magnetic pressure is negligible, then the system is unstable if $\\gamma_g -1 + \\alpha(\\gamma_{cr} -1) < 0$ \\citep{ZK75}. Thus, if the cosmic rays are well coupled by scattering and $\\gamma_{cr}=4/3$, they can stabilize the system, but if they diffuse rapidly, they are destabilizing. While a detailed investigation of this problem is beyond the scope of this paper, it is safe to say that if the instability occurs, its nonlinear development differs from that of the instability in a strong magnetic field.  If cosmic rays can be accelerated and confined in the presence of magnetic fields which are several orders of magnitude weaker than the fields in contemporary galaxies then evidence for cosmic rays in young galaxies is not evidence for magnetic fields at the equipartition level. The fieldstrengths required are several orders of magnitude larger than the seed fields produced by mechanisms operating on large scales, but are representative of the fields expected from a superposition of many small sources, such as plerion supernova remnants. The dominance of small scale fields predicted by some dynamo theories is inconsistent with the presence of cosmic rays, unless the field is strong enough that the cosmic ray gyroradius is less than the size of the system. Thus, the lower bounds on $B$ implied by the presence of cosmic rays impose meaningful constraints on theories of the origin and evolution of galactic magnetic fields (\\S\\ref{sec:early}).  Finally, the drastic changes in the properties of hydromagnetic turbulence brought about by cosmic rays may be of some importance beyond the present application, although they fall outside its immediate domain. It has been pointed out elsewhere that cosmic rays can be a significant source of hydromagnetic fluctuations at the scale of their gyroradius, and that this may enhance the rate at which they are accelerated in shocks [\\citet{E85}, \\citet{LB00}, \\citet{BL01}]. Because the gyroradius scale falls between the global galactic scale and the resistive scale, fluctuations driven by cosmic rays may affect the operation of the galactic dynamo. The novel dispersion relation of these fluctuations will affect their nonlinear behavior, and will modify the properties of a turbulent cascade. These effects could operate in any environment in which the magnetic field is weak and energetic particles are present."
    },
    "0212/astro-ph0212135_arXiv.txt": {
        "abstract": "We present observations and models of the behaviour of the HI and HeI lines between 1.6 and 2.2$\\mu$m in a small sample of compact HII regions.  As in our previous papers on planetary nebulae, we find that the `pure' 1.7007$\\mu$m 4$^{3}$D--3$^{3}$P and 2.16475$\\mu$m 7$^{3,1}$G--4$^{3,1}$F HeI recombination lines behave approximately as expected as the effective temperature of the central exciting star(s) increases.  However, the 2.058$\\mu$m 2$^1$P--2$^1$S HeI line does not behave as the model predicts, or as seen in planetary nebulae.  Both models and planetary nebulae showed a decrease in the HeI 2$^1$P--2$^1$S/HI Br$\\gamma$ ratio above an effective temperature of 40000K. The compact HII regions do not show any such decrease.  The problem with this line ratio is probably due to the fact that the photoionisation model does not account correctly for the high densities seen in these HII regions, and that we are therefore seeing more collisional excitation of the 2$^1$P level than the model predicts.  It may also reflect some deeper problem in the assumed model stellar atmospheres.  In any event, although the normal HeI recombination lines can be used to place constraints on the temperature of the hottest star present, the HeI 2$^1$P--2$^1$S/HI Br$\\gamma$ ratio should not be used for this purpose in either Galactic HII regions or in starburst galaxies, and conclusions from previous work using this ratio should be regarded with extreme caution.  We also show that the combination of the near infrared `pure' recombination line ratios with mid-infrared forbidden line data provides a good discriminant of the form of the far ultraviolet spectral energy distribution of the exciting star(s).  From this we conclude that CoStar models are a poor match to the available data for our sources, though the more recent WM-basic models are a better fit. ",
        "introduction": "This is the last paper in a series discussing the near infrared hydrogen and helium recombination lines in photoionised nebulae.  In the previous two papers (Lumsden, Puxley \\& Hoare 2001a,b; hereafter Paper 1 and 2) we discussed the properties as measured in planetary nebulae.  We showed that the ratio of pure hydrogen and helium recombination lines are a good measure of the hardness of the exciting radiation field from the central star.  This is as expected from simple photoionisation models, since the volume of the He$^+$ zone relative to the volume of the H$^+$ zone is simply related to the hardness of the radiation field in the 504--912\\AA\\ region of the spectrum of the exciting star.  In turn this depends on the stellar effective temperature as well as more specific details of the stellar spectrum as related to stellar evolution (ie stellar winds, luminosity, surface gravity etc).  For young stars, this ratio is also therefore clearly a constraint on the upper end of the stellar initial mass function (IMF).  The use of near infrared hydrogen and helium lines as probes of the underlying radiation field was suggested by Thompson \\& Tokunaga (1980) in the study of heavily embedded compact HII regions.  With early infrared instruments the near infrared pure recombination HeI lines were too weak to be useful.  Considerable attention was given instead to the HeI 2$^1$P--2$^1$S line at 2.058$\\mu$m. Unfortunately this line is largely pumped by the resonance in the HeI 2$^1$P--1$^1$S transition, and can be influenced by collisional excitation from the metastable 2$^1$S level as well as from the triplet series.  At first it was believed that the complexity of this transition could be modelled in a simple fashion.  Doyon, Puxley \\& Joseph (1992) suggested that the HeI 2$^1$P--2$^1$S to HI Br$\\gamma$ ratio was ideal for studying the initial mass function in starburst galaxies.  They tried to take account of the competing factors that led to emission in the 2.058$\\mu$m line using an analytical approach.  However, Shields (1993) showed from a full photoionisation treatment that this approach was a poor approximation to the actual predicted line strength.  More recent updates to the atomic data used in the prediction of the HeI 2$^1$P--2$^1$S line strength have also been presented by Ferland (1999).  In the previous two papers in this series we were able to show that the observed near infrared HeI and HI line strengths did largely follow the predictions of the photoionisation models for planetary nebulae.  This is encouraging since the planetary nebulae are at least simple in the sense that the objects we observed only had a single central exciting star, and could be modelled reasonably as expanding spheres.  If the model predictions for such simple systems deviated significantly from the observations, then the method of using near infrared HeI and HI line strengths as a constraint on the hardness of the stellar radiation field would have been invalidated almost completely.  The aim of this paper is to study how well the line ratios compare between model and observation for compact HII regions.  HII regions provide a more challenging test of these models for several reasons.  They probably contain a central star cluster rather than a single star, so the cumulative effect of the stellar output needs to be estimated. This in turn requires knowledge of the IMF.  Of course, this is the parameter that Doyon et al.\\ (1992) sought to estimate from the HeI to HI line ratios. However, in our case, where we are testing the models, we are forced to make some simple assumptions regarding the form of the initial mass function. Compact HII regions are also rather more diverse morphologically than the simple planetary nebulae we were considering.  Finally, there is considerable uncertainty about the form of the stellar radiation field in massive stars. Models diverge markedly depending on whether line blanketing due to metals is included or not, whether the effects of a stellar wind are included, and whether a non-LTE treatment is applied.  Martins, Schaerer \\& Hillier (2002) give a useful summary of the properties of the currently available models for dwarfs.  In addition to this overall uncertainty, there is also some doubt as to whether very young massive stars have the same effective radiation field as they would if they were on the main sequence.  We know that massive young stellar objects tend to have very dense, slow moving, stellar winds (eg.\\ Bunn, Drew \\& Hoare 1995).  These winds may reprocess much of the ionising radiation from the young stars into non-ionising radiation.  The objects studied by Bunn et al.\\ do not have substantial HII regions despite having the bolometric luminosity equivalent to OB stars.  The aim of this paper therefore is to determine how well we can reproduce the HeI/HI line ratios used in Papers 1 and 2 using realistic models of compact HII regions. ",
        "conclusions": "We have presented models for the infrared helium and hydrogen recombination line ratios in compact HII regions.  We fail to find any model that can accurately reproduce the HeI 2$^1$P--2$^1$S line strength as a function of effective temperature.  We therefore would discourage anyone from using this ratio as a constraint on the IMF, and any evidence derived from the literature about the IMF in a particular system that relies on this ratio should largely be discounted.  Our results do show reasonable agreement between model and observation for pure recombination lines such as HeI 4$^{3}$D--3$^{3}$P or 7$^{3,1}$G--4$^{3,1}$F when ratioed with HI Br$\\gamma$ as long as we use dwarf CoStar or WM-basic models.  This indicates that these are in principle a useful measure of the stellar effective temperature of the hottest star present.  The main drawback with these ratios is the relatively small range of effective temperature over which they are useful (approximately over the range B0 to O8 for main sequence CoStar models for example).  However, at least some of the galaxies in which the 1.7007$\\mu$m HeI 4$^{3}$D--3$^{3}$P line has been detected do have values of the ratio within this range (see, eg, Table 4 of Vanzi \\& Rieke 1997).  It should certainly help to address the issue of whether there is evidence for an anomalous IMF from infrared data (see, eg, the review by Scalo 1998).  We have outlined possible reasons for the discrepancies seen between the observations and the models for the HeI 2$^1$P--2$^1$S/HI Br$\\gamma$ ratio. The first, and perhaps most likely, is that the relative importance of photoionisation and collisional excitation from the $2^3$S HeI state may be incorrectly estimated by our HII region models.  Clearly, from Figure 9, models in which the collisional depopulation of $2^3$S is dominant come closer to matching the observed data.  However, it is still clear from Figure 12 that the HeI/HeI ratios are closer to a flat function of effective temperature than the declining function predicted by Cloudy (in agreement with the observations of the planetary nebulae in Papers 1 and 2).  It seems likely that errors in the 2$^3$S de-population rate may give rise to part but not all of the discrepancy seen.  Perhaps, instead, the HeI 2$^1$P--2$^1$S line is telling us something about the nature of the ultraviolet spectral energy distribution of the exciting star? The current uncertainty in the model stellar atmospheres for massive stars (eg Martins et al.\\ 2002) is unlikely to resolve this problem entirely however, since we have shown that the HeI/HI line ratios have a relatively weak dependence on model when comparing CoStar and WM-basic dwarf models.  The effective reduction in luminosity given by any re-calibration of the spectral type/effective temperature relation to cooler temperatures may help to lessen the differences between the model HeI 2$^1$P--2$^1$S/HI Br$\\gamma$ ratio and the observed data however.  Models with a more realistic treatment of the stellar wind (such as the WM-basic set) also help to reduce the discrepancy between the pure HeI/HI recombination ratios and the mid-infrared forbidden line ratios, since the overall shape of the far ultraviolet continuum is clearly crucial in that regard.  They also help in closing the gap between the direct spectral type seen in G29.96--0.02 and the inferred type from the photoionisation models. Similarly, consideration of changes in metallicity of the stellar model do not effect the HeI/HI ratios significantly, but can effect forbidden line ratios (eg Figure 5).  There is evidence for a gradient of metallicity with Galactocentric radius, with metallicity increasing inwards and excitation correspondingly decreasing (eg Giveon et al.\\ 2002).  Many of our sources lie within the solar circle and are likely to have exciting stars with higher than solar metallicity.  The most discrepant data point for all three helium lines is probably the youngest of the HII regions in the sample used here.  We have outlined the possibility for this case that the exciting star may actually resemble those massive young stellar objects studied by Bunn et al.\\ (1995).  In this case a considerable fraction of the emitted ultraviolet flux is actually reprocessed into non-ionising radiation by a dense stellar wind reducing the excitation in the nebular gas.  For the moment this must remain purely speculative however.  The easiest way of testing whether our results indicate some residual problems within the photoionisation model that are not otherwise apparent from Papers 1 and 2 or a deficiency on the assumed stellar models would be to obtain $K$ band spectra of older, more evolved HII regions, such as those studied by Kennicutt et al.\\ (2000).  More evolved regions are less dense, and larger, and hence any residual uncertainty in the depopulation route from 2$^3$S should be removed. If the same effect as seen here were repeated then it would tend to indicate that even the latest stellar models are still not well matched to the actual ultraviolet continuum of OB stars."
    },
    "0212/astro-ph0212303_arXiv.txt": {
        "abstract": "We have carried out an investigation of the abundance of deuterium along two extended sight lines through the interstellar medium (ISM) of the Galactic disk. The data include {\\it Far Ultraviolet Spectroscopic Explorer (FUSE)} observations of HD~195965 (B1Ib) and HD~191877 (B0V), as well as Space Telescope Imaging Spectrograph (STIS) observations of HD~195965. The distances to HD~195965 and HD~191877, derived from spectroscopic parallax, are $794\\pm200$~pc and $2200\\pm550$~pc, respectively, making these the longest Galactic disk sight lines in which deuterium has been investigated with \\fuse. The \\fuse\\ spectra contain all of the \\ion{H}{1} Lyman series transitions (and the corresponding D transitions) except Ly$\\alpha$. The higher Lyman lines clearly show the presence of deuterium. We use a combination of curve of growth analyses and line profile fitting to determine the \\ion{D}{1} abundance toward each object.  We also present column densities for \\ion{O}{1} and \\ion{N}{1} toward both stars, and \\ion{H}{1} measured from Ly$\\alpha$ absorption in the STIS spectrum of HD~195965. Toward HD~195965 we find D/H=$(0.85\\pm^{0.34}_{0.24})\\times10^{-5}$ ($2\\sigma$), O/H=$(6.61\\pm^{1.03}_{1.11})\\times10^{-4}$, and N/H=$(7.94\\pm^{1.69}_{1.34})\\times10^{-5}$. Toward HD~191877 we find D/H=$(0.78\\pm^{0.52}_{0.25})\\times10^{-5}$ ($2\\sigma$) and N/H=$(6.76\\pm^{2.22}_{1.97})\\times10^{-5}$. The \\ion{O}{1} column density toward HD~191877 is very uncertain. Our preferred value gives O/H=$(3.09\\pm^{1.98}_{0.98})\\times10^{-4}$, but we cannot rule out O/H values as low as O/H=$1.86\\times10^{-4}$, so the O/H value for this sight line should be taken with caution. The D/H ratios along these sight lines are lower than the average value of ($1.52\\pm0.15)\\times10^{-5}$ ($2\\sigma$ in the mean) found with \\fuse\\ for the local interstellar medium ($\\sim 37$ to 179 pc from the Sun). These observations lend support to earlier detections of variation in D/H over distances greater than a few hundred pc. The O/H ratio toward HD~195965 is supersolar. This star is part of an OB association, so there may be local enrichment by nearby massive stars. The D/H and O/H values measured along these sight lines support the expectation that the ISM is not well mixed on distances of $\\sim1000$ pc. These observations demonstrate that although D/H studies through Lyman absorption may become impractical at d$>$2500 pc and log $N$(\\ion{H}{1})$>$21, D/H studies in the distance range from 500 to 2500 pc may be very useful for investigating mixing and chemical evolution in the ISM. ",
        "introduction": "The abundance of deuterium in the interstellar medium (ISM) is a strong constraint on chemical evolution scenarios. Unlike most elements, the net cosmic abundance of deuterium is not enhanced by stellar processes or supernovae, so its abundance has been declining from the primordial value due to destruction by stellar processing. The primordial abundance of deuterium is itself an important quantity, as it can be used to measure $\\eta$, the universal ratio of baryons to photons, which can be compared to the predictions of Big Bang nucleosynthesis \\citep{rafs73,bs85,bnt01}.  One approach to finding the initial deuterium abundance is to measure it in high-redshift QSO absorption-line systems, assuming that the clouds which give rise to these systems have undergone very little stellar processing \\citep{w97,tblfws99,otkspl01,pb01,ldd02}. However, the large range in D/H values measured in these systems suggests that the gas in some of these clouds may in fact have undergone some stellar processing \\citep{pi01,foscv01}. It is therefore crucial to understand how chemical evolution affects the abundance of deuterium relative to other elements. Several groups are measuring the abundance of deuterium and other elements along many sight lines in the Galactic ISM in order to establish the mean D/H ratio in the Galactic disk, and to look for variations in the D/H ratio. These variations can shed light on the effects of star-formation history and mixing processes on the D/H ratio. The local D/H value can then be compared to regions that have undergone varying amounts of stellar processing ({\\it e.g.,} high velocity clouds, or intergalactic \\ion{H}{1} clouds), which will illuminate the effects of stellar processing on the D/H ratio in these systems.  The \\copernicus\\ satellite was used to measure the \\ion{D}{1} and \\ion{H}{1} Lyman series absorption toward $\\sim20$ stars, and the results suggested a D/H ratio of 1.5$\\times$10$^{-5}$, with a dispersion that may reflect real variations in the ISM (see Vidal-Madjar \\& Gry 1984 for a review). Several measurements of D/H in the local ISM have been made with the {\\em Hubble Space Telescope} (\\hst) \\citep{lin95,l96,lin02,lem96,v98,h99,s99,vpltcg00}, though only the Ly$\\alpha$ transitions of \\ion{D}{1} and \\ion{H}{1} could be analyzed, limiting the number of available sight lines to those with low enough \\ion{H}{1} columns that the \\ion{D}{1} absorption is not blended with the \\ion{H}{1} line but large enough \\ion{H}{1} columns that \\ion{D}{1} can be detected. Three sight lines observed with the Interstellar Medium Absorption Profile Spectrograph (\\imaps) spanned larger distances than the \\hst\\ sight lines (250 to 500 pc versus $\\la$ 100 pc), and like the \\copernicus\\ results they showed evidence of variation in D/H with location \\citep{j99,stfjvw00}.  One of the science goals of the {\\em Far Ultraviolet Spectroscopic Explorer} (\\fuse) is to measure D/H in the Milky Way \\citep{m00}. \\fuse\\ is uniquely suited to do this because it observes the spectral region from 905 to 1187~\\AA\\, which includes all of the \\ion{D}{1} Lyman transitions except Ly$\\alpha$. The ability to use multiple \\ion{D}{1} transitions means that \\fuse\\ is able to measure $N$(\\ion{D}{1}) along higher column density sight lines than missions that only observe Ly$\\alpha$ ({\\it e.g.,} \\hst), and for many more sight lines than short term missions ({\\it e.g.,} \\imaps). Moos et al. (2002, and references therein) reported results for the first set of seven sight lines ranging from 37 to 179 pc, five of which had both \\ion{D}{1} measurements (from \\fuse) and \\ion{H}{1} (from \\hst). The results were consistent with a single value of D/H$=(1.52\\pm0.08)\\times10^{-5}$ (the uncertainty is $1\\sigma$ in the mean). However, none of these sight lines extend as far as the three sight lines observed by \\imaps. Analysis of more extended sight lines is crucial for the verification and understanding of variations in the D/H ratio in the ISM.  In this paper we present \\fuse\\ measurements of D/H toward two distant stars in the Galactic disk: HD~195965 and HD~191877. These stars are at distances greater than the \\imaps\\ targets, extending the distance and range of $N$(\\ion{H}{1}) for D/H measurements in the ISM. Information about these two sight lines is listed in Table 1. The \\fuse, STIS, and ground-based observations and data reduction are described in \\S\\ 2. Section 3 describes some properties of the sight lines. In \\S\\ 4 we describe both the curve of growth and the absorption profile fitting analyses of the \\ion{D}{1} column density toward both stars. In \\S\\ 5 we discuss the determination of the \\ion{O}{1} and \\ion{N}{1} abundances along these sight lines, and in \\S\\ 6 we re-examine the \\ion{H}{1} column densities. Section 7 contains a discussion of the results and conclusions. ",
        "conclusions": "The sight lines analyzed in this paper are the longest ones for which the \\ion{D}{1} column densities have been determined with \\fuse. As such, they present the opportunity to look for variations in \\ion{D}{1}/\\ion{H}{1}, \\ion{D}{1}/\\ion{O}{1}, and \\ion{D}{1}/\\ion{N}{1} beyond the local ISM sight lines presented by \\cite{m02}. Table 8 contains column density ratios of various elements toward both stars, along with the mean values for the local ISM from \\cite{m02}. In this section we discuss the behavior of these column density ratios beyond the local ISM.  \\subsection{D/H}  The deuterium abundances along seven sight lines have been analyzed with \\fuse\\ \\citep{m02,k02,f02,s02,l02,nl02,w02,h02}. Five of these had both \\ion{D}{1} and \\ion{H}{1} measurements, and the results were consistent with a constant value of D/H in the local ISM. Earlier \\copernicus\\ observations showed evidence for variations in D/H at larger distances (see Vidal-Madjar \\& Gry 1984 for a review), evidence which was strengthened by subsequent \\imaps\\ investigations \\citep{j99,stfjvw00}. The main obstacle to verifying and understanding this variation is the paucity of suitable deuterium sight lines that probe distances $\\ge$250 pc.  Figure 16 shows the D/H ratios for HD~195965 and HD~191877 plotted as a function of distance from the Sun, along with five of the seven original \\fuse\\ sight lines with reliable \\ion{H}{1} measurements, and several others from \\hst, \\imaps, and \\copernicus\\ (see Moos et al. 2002).  The D/H ratios toward both stars are lower than the average value found for the original seven \\fuse\\ sight lines. This contrasts with the lack of variation seen in the local ISM. The new sight lines support the existence of variations in the D/H ratio at distances larger than a few hundred pc, as also suggested by \\imaps\\ and\\copernicus.  In the simplest model of chemical evolution, stellar processing (astration) leads to a decline in the D/H ratio because D is easily destroyed in stars. Thus, the D/H ratio is expected to vary among regions of the Galaxy that have undergone different amounts of star formation in the past, provided that the gas in these regions has not mixed ({\\it e.g.} Tosi et al. 1998; Chiappini et al. 2002). In a small volume like the local ISM ($r=100$ pc), \\cite{m02} showed that the time between supernova $t_{SN}$ is roughly equal to the mixing time scale (assumed to be the sound crossing time $t_{s}$), suggesting that the gas in such a region is well-mixed. If a larger region is considered, say 1000 pc in radius to approximate the scales probed by HD~195965 and HD~191877, the crossing time increases by a factor of 10, while the time between supernovae decreases by a factor of 100, so that $t_{SN} \\ll t_{s}$. Supernovae are not distributed randomly over this region but are spatially correlated in regions of star formation, meaning that substantial inhomogeneities will arise between star-forming regions and more quiescent regions. For these reasons it is not surprising to see variations over the distances probed by HD~195965 and HD~191877.  A more precise treatment of mixing in the ISM was carried out by \\cite{am02}. Using numerical simulations, they found that the mixing time scales were a few times 10$^8$ years. Chemical inhomogeneities in the ISM are long-lived, so variations in D/H are expected over 1000~pc length scales. In fact, since the mixing time scales increase only slightly with length scale in their models, it is unlikely that the local ISM ({\\it e.g.,} the original seven \\fuse\\ sight lines) would have a constant D/H value if any sources of inhomogeneities ({\\it i.e.,} supernovae) have occurred within the past $\\sim10^8$ years.  We have investigated several alternative explanations for the low D/H values. One possibility is that some of the D column is in high-velocity components which are outside of the velocity integration range used to determine $N$(\\ion{D}{1}) \\citep{j99}. We have used the \\ion{O}{1} 1302.169~\\AA\\ line in the STIS spectrum of HD~195965 to rule this out for that sight line (see \\S\\ 5). Note that $\\sim50\\%$ of the \\ion{D}{1} column would have to be missed if the D/H ratio for the HD~191877 sight line were as high as the \\cite{m02} mean local ISM value. Figure 7 shows that this is not the case, although we cannot rule out a much smaller systematic offset for HD~191877.  A second possibility is that a significant amount of D is locked up in HD molecules. To test this idea, we have measured the equivalent widths of several strong, isolated Lyman series HD lines in the \\fuse\\ spectra: (3-0)~R(0)~1066.271~\\AA, (5-0)~R(0)~1042.847~\\AA, and (4-0)~R(0)~1054.288~\\AA. The column densities derived from the equivalent widths are $\\sim2.4\\times10^{14}$ cm$^{-2}$ for the HD~195965 sight line, and $\\sim3.0\\times10^{14}$ cm$^{-2}$ for the HD~191877 sight line. For both sight lines this is $\\la3.5$\\% of the total D column density, much lower than the uncertainty in $N$(\\ion{D}{1}).  \\subsection{O/H}  The D/O and O/H ratios toward HD~195965 in Table 8 are quite different from the average values in the local ISM found by \\cite{m02}. This is a surprising result, since \\cite{mjc98} found very little variation in O/H in their survey, which covered distances and \\ion{H}{1} column densities similar to those of the sight lines in this paper.  Both the \\ion{O}{1} and \\ion{H}{1} column densities toward HD~195965 were determined through at least two independent methods, with very similar results. The \\ion{H}{1} was measured using both \\iue\\ and STIS data, and the \\ion{O}{1} was measured with a single-component COG which combined \\fuse\\ and STIS data, as well as independent estimates from the optically thin 1355.598~\\AA\\ line and the damped 1302.169~\\AA\\ line in the STIS spectrum. Both of these column densities are therefore quite firmly established, and we believe that it is very unlikely that there is a substantial systematic error in either quantity.  HD~195965 appears to lie on a very high metallicity sight line, with supersolar O/H \\citep{sm01}. This may be a region where the signature of stellar processed gas is still visible (as is not the case for the local bubble). \\cite{hum78} found that this star is in the Cygnus OB7 association. If so, it may be in a peculiar region, perhaps in a metal-enriched cloud of ejecta from the star itself or a neighbor. Since gas with high metallicity is thought to have undergone stellar processing, D/H in such gas is expected to be low. This expected anticorrelation between O/H (or N/H) and D/H was not seen in the first seven \\fuse\\ sight lines, although it could have been masked by the small number of sight lines studied \\citep{st02}. Figure 17 compares the O/H and D/H value for HD~195965 with those of the original seven \\fuse\\ sight lines. Unfortunately HD~195965 is the only point at high metallicity, so no firm conclusions can be drawn. However, it is possible that this sight line shows evidence of the expected anticorrelation in the ISM (see also Steigman 2002). Note that an error in the \\ion{H}{1} column density toward HD~195965 would tend to move that point along a diagonal line with a positive slope, and thus could not produce the expected trend of decreasing D/H with rising O/H.  The HD~191877 sight line lacks STIS data, so we could not constrain the \\ion{O}{1} and \\ion{H}{1} abundances with multiple estimates. For these reasons we are less confident in the \\ion{O}{1} column densities found for the HD~191877 sight line. As discussed in \\S\\ 5, the derived log $N$(\\ion{O}{1}) is very uncertain, although our preferred value, found by coupling the $b$-values of \\ion{N}{1} and \\ion{O}{1}, is consistent with the \\fuse\\ local ISM ratio and the \\cite{mjc98} ratio. Unfortunately, with the existing information, we cannot make any strong conclusions about O/H toward HD~191877.  \\subsection{N/H}  The N/H ratios toward HD~195965 and HD~191877 in Table 8 are higher than the average values in the local ISM found by \\cite{m02}. Figure 18 shows a plot of N/H versus D/H along both sight lines, compared to the four \\fuse\\ sight lines for which N/H was measured.  The figure is very suggestive of an anticorrelation between D/H and N/H.  Complicating the interpretation of this plot is the fact N/H ratio is susceptible to ionization effects, especially at low column densities which do not provide adequate shielded from ionizing photons, such as the column densities along the local ISM sight lines \\citep{j00}. The dashed line in Figure 18 shows the ISM value of N/H found by \\cite{mcs97}. The sight lines used in that study had higher \\ion{H}{1} column densities than the \\fuse\\ local ISM sight lines, so the \\fuse\\ sight lines are very likely underestimating N/H, because a significant fraction of N is ionized. The apparent trend in Figure 18 may simply reflect the tendency for $N$(\\ion{N}{2})/$N$(\\ion{N}{1}) to decrease with increasing $N$(\\ion{H}{1}).  However, an anticorrelation between D and N abundances would be expected if D is destroyed in stars while N is produced ({\\it e.g.} Steigman et al. 2002). The overabundance of heavy elements accompanied by the underabundance of D seems to fit well within a simple chemical evolution scenario. While ionization effects probably play a role in the high N/H ratio on these sight lines, they may not be the only factor. A comprehensive N/H survey using \\fuse\\ spectra is currently underway, and may shed light on this issue.  Several sight lines do not follow the expected trend of low D/H accompanied by high N/H or O/H. \\cite{s02} found a low O/H ratio toward BD+28$^{\\circ}$4211 with \\fuse, while their D/H ratio is close to the Local ISM mean. \\cite{j00} found a low D/H toward $\\delta$ Orionis using \\imaps, but also a somewhat low N/H ratio. These deviations suggest that the simple picture of increasing heavy elements abundances with decreasing D/H is not complete.  \\subsection{Future Prospects for D/H Measurements}  Measuring D/H values in the Milky Way becomes increasingly difficult as $N$(\\ion{H}{1}) increases, perhaps reaching a fundamental confusion limit. The COGs in Figures 8 and 9 show that the high order \\ion{D}{1} Lyman transitions are close to saturation. At higher column densities, these transitions will fall on the flat part of the COG and will not give precise column density information. Even higher order transitions at shorter wavelengths, where the \\ion{D}{1} lines might not yet be saturated at higher column density, are not useful because the \\ion{H}{1} transitions begin to blend together. These facts make it unlikely that many Galactic disk sight lines more distant than those studied here will be useful for D/H measurements. However, the sight lines analyzed in this paper demonstrate that D/H properties over the distance range from 500 to 2500 pc can be explored with \\fuse, even if these sight lines contain a significant amount of H$_2$. This first glimpse at D/H in this distance range suggests that it will be very useful for understanding mixing and chemical evolution in the ISM.  The chemical evolution model of \\cite{crm02} predicts an increase in D/H with galactocentric radius, accompanied by a decrease in N/H and O/H, leading to a pronounced gradient in D/O and N/O. The gas in the outer parts of the galaxy has undergone less stellar processing, and in fact \\cite{crm02} predict that the D/H ratios in the outer disk will approach the primordial value. The sight lines studied here have roughly the same galactocentric radii as the local ISM sight lines of \\cite{m02}, so we cannot test this model, but with \\fuse\\ observations of sight lines in the inner and outer Galactic disk this may prove possible."
    }
}